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Patrick Moore’s Practical Astronomy Series 


Other Titles in this Series 

Navigating the Night Sky 
How to Identify the Stars and 
Constellations 
Guilherme de Almeida 

Observing and Measuring Visual 
Double Stars 
Bob Argyle (Ed.) 

Observing Meteors, Comets, Supernovae 
and other transient Phenomena 
Neil Bone 

Human Vision and The Night Sky 
How to Improve Your Observing Skills 
Michael P. Borgia 

How to Photograph the Moon and Planets 
with Your Digital Camera 
Tony Buick 

Practical Astrophotography 
Jeffrey R. Charles 

Pattern Asterisms 
A New Way to Chart the Stars 
John Chiravalle 

Deep Sky Observing 
The Astronomical Tourist 
Steve R. Coe 

Visual Astronomy in the Suburbs 
A Guide to Spectacular Viewing 
Antony Cooke 

Visual Astronomy Under Dark Skies 
A New Approach to Observing Deep Space 
Antony Cooke 

Real Astronomy with Small Telescopes 
Step-by-Step Activities for Discovery 
Michael K. Gainer 

The Practical Astronomer’s Deep-sky 

Companion 

Jess K. Gilmour 

Observing Variable Stars 
Gerry A. Good 

Observer’s Guide to Stellar Evolution 
The Birth, Life and Death of Stars 
Mike Inglis 

Field Guide to the Deep Sky Objects 
Mike Inglis 


Astronomy of the Milky Way 
The Observer’s Guide to the 
Southern/Northern Sky Parts 1 and 2 
hardcover set 
Mike Inglis 

Astronomy of the Milky Way 
Part 1: Observer’s Guide to the 
Northern Sky 
Mike Inglis 

Astronomy of the Milky Way 
Part 2: Observer’s Guide to the 
Southern Sky 
Mike Inglis 

Observing Comets 

Nick James and Gerald North 

Telescopes and Techniques 
An Introduction to Practical Astronomy 
Chris Kitchin 

Seeing Stars 

The Night Sky Through Small Telescopes 
Chris Kitchin and Robert W. Forrest 

Photo-guide to the Constellations 
A Self-Teaching Guide to Finding Your 
Way Around the Heavens 
Chris Kitchin 

Solar Observing Techniques 
Chris Kitchin 

How to Observe the Sun Safely 
Lee Macdonald 

The Sun in Eclipse 

Sir Patrick Moore and Michael Maunder 
Transit 

When Planets Cross the Sun 

Sir Patrick Moore and Michael Maunder 

Light Pollution 
Responses and Remedies 
Bob Mizon 

Astronomical Equipment for Amateurs 
Martin Mobberley 

The New Amateur Astronomer 
Martin Mobberley 

Lunar and Planetary Webcam User’s Guide 
Martin Mobberley 


(Continued after Index) 


Astrophysics 
is Easy! 

An Introduction 
for the Amateur 
Astronomer 


Mike Inglis 


Springer 


Dr Mike Inglis FRAS 
SUNY 

inglism@sunysuffolk.edu 


Library of Congress Control Number: 2007925262 

Apart from any fair dealing for the purposes of research or private study, or criticism or review, 
as permitted under the Copyright, Designs and Patents Act 1988, this publication may only be 
reproduced, stored or transmitted, in any form or by any means, with the prior permission in writing 
of the publishers, or in the case of reprographic reproduction in accordance with the terms of licences 
issued by the Copyright Licensing Agency. Enquiries concerning reproduction outside those terms 
should be sent to the publishers. 

Patrick Moore’s Practical Astronomy Series ISSN 1617-7185 
ISBN-13: 978-1-85233-890-9 e-ISBN-13: 978-1-84628-736-7 

Springer Science+Business Media 
Springeronline.com 

©Springer- Verlag London Limited 2007 

The use of registered names, trademarks, etc. in the publication does not imply, even in the absence 
of a specific statement, that such names are exempt from the relevant laws and regulations and 
therefore free for general use. 

The publisher makes no representation, express or implied, with regard to the accuracy of the 
information contained in this book and cannot accept any legal responsibility or liability for any 
errors or omissions that may be made. Observing the Sun, along with a few other aspects of 
astronomy, can be dangerous. Neither the publisher nor the author accepts any legal responsibility 
or liability for personal loss or injury caused, or alleged to have been caused, by any information or 
recommendation contained in this book. 


For Dad and Alan, who are already amongst the stars 


Contents 


Preface and Thanks xi 

Acknowledgements xiii 

Overview xv 

Chapter 1 Tools of the Trade 1 

1.1 Distance 1 

1.2 Brightness and Luminosity 6 

1.3 Magnitudes 8 

1.4 Color 15 

1.5 Size and Mass 19 

1.6 Star Constituents 22 

1.7 Spectra and Spectroscopy 23 

1.8 Stellar Classification 25 

1.9 The Hertzsprung- Russell 

Diagram 35 

1.10 The H-R Diagram and Stellar Radius 37 

1.11 The H-R Diagram and Stellar Luminosity 39 

1.12 The H-R Diagram and Stellar Mass 39 

Chapter 2 The Interstellar Medium 45 

2.1 Introduction 45 

2.2 Nebulae 47 


VII 


VIII 


Contents 


2.3 Emission Nebulae 47 

2.4 Dark Nebulae 53 

2.5 Reflection Nebulae 56 

2.6 Molecular Clouds 57 

2.7 Protostars 58 

2.8 The Jeans Criterion 59 

Chapter 3 Stars 63 

3.1 The Birth of a Star 63 

3.2 Pre-Main-Sequence Evolution and the Effect 

of Mass 66 

3.3 Mass Loss and Gain 70 

3.4 Clusters and Groups of Stars 72 

3.5 Star Formation Triggers 84 

3.6 The Sun — The Nearest Star 86 

3.7 Binary Stars and Stellar Mass 92 

3.8 Lifetimes of Main-Sequence Stars 97 

3.9 Red Giant Stars 101 

3.10 Helium-Burning and the Helium Flash 104 

3.11 Star Clusters, Red Giants, and the H-R Diagram .... 107 

3.12 Post-Main-Sequence Star Clusters: The Globular 

Clusters 108 

3.13 Pulsating Stars 114 

3.14 The Death of Stars 122 

3.15 The Asymptotic Giant Branch 122 

3.16 Dredge-Ups 124 

3.17 Mass Loss and Stellar Winds 125 

3.18 Infrared Stars 125 

3.19 The End of an AGB Star’s Life 126 

3.20 Planetary Nebulae 128 

3.21 White Dwarf Stars 133 

3.22 High-Mass Stars and Nuclear Burning 138 

3.23 Iron, Supernovae, and the Formation 

of the Elements 141 

3.24 The End Result of High-Mass Stars’ Evolution: 

Pulsars, Neutron Stars, and Black Holes 147 

Chapter 4 Galaxies 157 

4.1 Introduction 157 

4.2 Galaxy Types 158 

4.3 Galaxy Structure 158 

4.4 Stellar Populations 159 

4.5 Hubble Classification of Galaxies 159 

4.6 Observing Galaxies 161 

4.7 Active Galaxies and AGNs 177 

4.8 Gravitational Lensing 182 

4.9 Redshift, Distance, 

and the Hubble Law 184 


Contents ix 

4.10 Clusters of Galaxies 185 

4.11 Endnote 188 

Appendix 1 Degeneracy 191 

Appendix 2 Books, Magazines, and Astronomical Organizations 193 

Books, Magazines, and Organizations 193 

Star Atlases and Observing Guides 193 

Astronomy and Astrophysics Books 194 

Magazines 195 

Organizations 195 

Topic Index 197 

Object Index 201 


Preface and 
Thanks 


Once again, I took paper to pen, and began a journey to explain the mysterious 
and beautiful complexities of stars, galaxies, and the material that lies between 
them. It was a journey that took many roads with many side-turnings as I 
often spent many long, lonely hours worrying whether I was being too obtuse, 
or at times patronizing. It is a fact that many amateur astronomers are very 
knowledgeable of the subject that they pursue with a passion. However, the book 
eventually came into sight, and this, for me a mammoth task, was completed. 
You now hold it in your hands! 

Throughout the entire process of writing the book, I was lucky enough to have 
the support of my publisher, Harry Blom, who, as a professional astronomer 
himself, knows only too well that astronomy authors are a breed apart and need 
to be pampered and dealt with extreme patience. Thanks, Harry — I owe you 
a pint. I must also thank my great friend John Watson, also associated with 
Springer, who gave the initial thumbs-up when I first outlined the idea for the 
book. John is an amateur astronomer himself, so he knows exactly what should 
go into a book, and perhaps even more importantly, what should be left out! I 
also owe you a pint. 

I am fortunate to have been taught astronomy by some of the world’s leading 
experts, and it was, and still is, a privilege to know them. In my humble 
opinion, not only are they superb astronomers, whether theoretical or observa- 
tional, but also wonderful educators. They are Chris Kitchin, Alan McCall, Iain 
Nicolson, Robert Forrest, and the late Lou Marsh. They were the best teachers I 
ever had. 


XI 


XII 


Preface and Thanks 


During the time spent writing this book, usually alone, usually at night, usually 
tired, I had the company of some wonderful musicians whose music is truly 
sublime. They are Steve Roach, David Sylvian, John Martyn, and the Blue Nile. 

Many friends have helped raise my spirits during those times when not all was 
going right according to the Inglis Master Plan. They listened to me complain, 
laughed at my jokes, and helped me remain sane — for the most part. So I want 
to say thank you to my British friends — Pete, Bill, Andy and Stuart — and my 
new friends here in the USA — Sean and Matt. It is nice to know that beer is the 
universal lubricant of friendship, whether it is McMullens or Blue Point. 

Astronomy is a very important part of my life, but not as important as my 
family; my brother, Bob, is a great friend and a strong source of support, especially 
during my formative years as an astronomer. My mother, Myra, is amazing, full 
of energy, spirit, and laughter, and has been supportive of my dream to be an 
astronomer since I was knee-high to a tripod. She is truly an example to us all. 
And of course Karen, my partner. I am not exaggerating when I say this book 
would not have seen the light of day without her help. “Diolch Cariad.” 

For making my life so much fun, cheers! 

Dr. Mike Inglis 
Long Island, USA, 2006 


Acknowledgements 


I would like to thank the following people and organisations for their help and 

permission to quote their work and for the use of the data they provided: 

The European Space Organization, for permission to use the Hipparcos and Tycho 
catalogues. 

My colleagues at Suffolk County Community College, USA, for their support and 
encouragement. 

The astronomers at Princeton University, USA, for many helpful discussions on 
the whole process of star formation. 

The astronomers at the University of Hertfordshire, UK, for inspirational lectures 
and discussions. 

Gary Walker, of the American Association of Variable Star Observers, for infor- 
mation on the many types of variable stars. 

Cheryl Gundy, of the Space Telescope Science Institute, USA, for supplying astro- 
physical data on many of the objects discussed. 

Dr. Stuart Young, of the University of Hertfordshire, UK, for discussions and 
information relating to star formation and the Hertzsprung-Russell diagram, 
and impromptu tutorials on many aspects of astronomy. 

Dr. Chris Packham, of the University of Florida, USA, for his help on pointing 
out several mistakes I have made over the years and for his input regarding 
AGNs. 

Karen Milstein, for the superb and professional work that she did reading through 
the initial proofs of the book, when there seemed to be more errors than facts! 


XIII 


xiv Acknowledgements 

The Smithsonian Astrophysical Observatory, USA, for providing data on many 
stars and star clusters. 

Robert Forrest, of the University of Hertfordshire Observatory, UK, for use of his 
observing notes. 

Michael Hurrell and Donald Tinkler of the South Bayfordbury Astronomical 
Society, UK, for use of their observing notes. 

In developing a book of this type, which presents a considerable amount of detail, 
it is nearly impossible to avoid error. If any arise, I apologize for the oversight, 
and I would be more than happy to hear from you. Also, if you feel that I have 
omitted a star or galaxy that you think would better describe a certain aspect of 
astrophysics, please feel free to contact me at: inglism@sunysuffolk.edu. I can’t 
promise to reply to all e-mails, but I will certainly read them. 


Overview 


To most normal people, astrophysics — the science of stars, galaxies, and the 
universe we live in — would seem to be a topic suited to a university-level 
textbook, and so the idea of a guide to astrophysics for the amateur astronomer 
may not, on first appearance, make any sense. However, let me assure you that 
anyone can understand how a star is born, lives its life, and dies, how galaxies 
are thought to evolve and what their shape can tell us about their origins and 
age, and even how the universe began and how it may end. In fact, very little 
mathematics is needed, and when it is used, it is only a matter of multiplication, 
division, subtraction, and addition! 1 

What is more, there are many wonderful objects that can be observed in the 
night sky that will illustrate even the most obtuse astrophysics concepts. All one 
needs is a willingness to learn and a dark night sky. 

Learning about, say, the processes that give rise to star formation, or what 
happens to a very large star as it dies, or even why some galaxies are spiral 
in shape whereas others are elliptical can add another level of enjoyment and 
wonder to an observing session. For instance, many amateur astronomers are 
familiar with the star Rigel, in the constellation Orion, but how many of you 
know that it is a giant star, with a mass more than 40 times that of the Sun, and 
it is nearly half a million times more luminous than the Sun? How many know 
that the closest large galaxy, M31 in Andromeda, has a supermassive black hole 
lurking at its center with a mass more than 50 million times that of the Sun? Or 
that the Orion Nebula, regarded by many as the premier nebula in the sky, is in 
fact an enormous stellar nursery where stars are actually being born as you read 
this book? Knowing details such as these can add another level of enjoyment to 
your observing sessions. 


XV 


XVI 


Astrophysics is Easy 


Each section of this book addresses a specific aspect of astrophysics. The 
first part focuses on the concepts needed for a complete understanding of the 
remainder of the book, and as such will be divided into specific topics, such as 
the brightness, mass, and distance of stars, and so on. Then we will look at the 
tools of an astronomer, namely spectroscopy. It is true to say that nearly all of 
what we know about stars and galaxies was and is determined by this important 
technique. We shall spend a fair amount of time looking at something called the 
Hertzsprung-Russell diagram; if ever a single concept or diagram could epitomize 
a star’s life (and even a star cluster’s life), the H-R diagram, as it is known, is the 
one to do it. It is perhaps the most important and useful concept in all of stellar 
evolution, and it is fair to say that once you understand the H-R diagram, you 
understand how a star evolves. 

Moving on to the objects themselves, we start with the formation of stars from 
dust and gas clouds, and conclude with the final aspect of a star’s life, which can 
end in a spectacular event known as a supernova, resulting in the formation of 
a neutron star and perhaps a black hole! 

On a grander scale, we delve into galaxies, their shapes (or morphology, as it 
is called), distribution in space, and origins. 

The topics covered are chosen specifically so that examples of objects under 
discussion can be observed; thus, at every point in our journey, an observing 
section will describe the objects that best demonstrate the topics discussed. Many 
of the objects, whether they be stars, nebulae, or galaxies, will be visible with 
modest optical instruments, and many with the naked eye. In a few exceptional 
cases, a medium-aperture telescope may be needed. Of course, not all observable 
objects will be presented, but just a representative few (usually the brightest 
examples). These examples will allow you to learn about stars, nebulae, and 
galaxies at your own pace, and they will provide a detailed panorama of the 
amazing objects that most of us observe on a clear night. 

For those of you who have a mathematical mind, some mathematics will be 
provided in the specially labelled areas. But, take heart and fear not — you do not 
have to understand any mathematics to be able to read and understand the book; 
it is only to highlight and further describe the mechanisms and principles of 
astrophysics. However, if you are comfortable with maths, then I recommend that 
you read these sections, as they will further your understanding of the various 
concepts and equip you to determine such parameters as a star’s age and lifetime, 
distance, mass, and brightness. All of the maths presented will be simple, of a 
level comparable to that of a 4th-year school student, or an 8th-grader. In fact, 
to make the mathematics simpler, we will use rough (but perfectly acceptable) 
approximations and perform back-of-the-envelope calculations, which, surpris- 
ingly, produce rather accurate answers! 

An astute reader will notice immediately that there are no star maps in the 
book. The reason for this is simple: in previous books that I have written, 
star maps were included, but their size generated some criticism. Some readers 
believed the maps were too small, and I tend to agree. To be able to offer large 
and detailed star maps of every object mentioned in this book would entail a 
doubling of its size, and probably a tripling of cost. With the plethora of star-map 
software that is available these days, it is far easier for readers to make their own 
maps than to present any here. 


Overview 


XVII 


A final point I wish to emphasize here is that the book can be read in several 
ways. Certainly, you can start at the beginning and read through to the end. 
But if you are particularly interested in, say, supernovae and the final stages 
of a star’s life, or in galaxy clusters, there is no reason that you should not go 
straight to that section. Some of the nomenclature might be unfamiliar, but I 
have attempted to write the book with enough description that this should not 
be a problem. Also, many of you will undoubtedly go straight to the observing 
lists. Read the book in the way that is most enjoyable to you. 

Without further ado, let us begin a voyage of discovery... 


Note 


1. Well, o.k. — we do use powers of ten occasionally, and numbers multiplied 

by themselves from time to time. But nothing else... honest! 




CHAPTER ONE 


Tools of the Trade 


1.1 Distance 


To determine many of the basic parameters of any object in the sky, it is first 
necessary to determine its proximity to us. We shall see later how this is vitally 
important because a star’s bright appearance in the night sky could signify that 
it is close to us or that it is an inherently bright star. Conversely, some stars may 
appear faint because they are at immense distance from us or because they are 
very faint stars in their own right. We need to be able to determine which is the 
correct explanation. 

Determining distance in astronomy has always been, and continues to be, 
fraught with difficulty and error. There is still no consensus as to which is the 
best method, at least for distances to other galaxies and to the farthest edges of 
our own galaxy — the Milky Way. The oldest method, still used today, is probably 
the most accurate, especially for determining the distances to stars. 

This simple technique is called Stellar Parallax. It is basically the angular 
measurement when the star is observed from two different locations on the 
Earth’s orbit. These two positions are generally six months apart, and so the star 
will appear to shift its position with respect to the more distant background stars. 
The parallax (p) of the star observed is equal to half the angle through which its 
apparent position appears to shift. The larger the parallax (p), the smaller the 
distance (d) to the star. Figure 1.1 illustrates this concept. 

If a star has a measured parallax of 1 arcsecond (l/3600th of a degree) 
and the baseline is 1 astronomical unit (AU), which is the average distance 


1 


2 


Astrophysics is Easy 



Earth 

[January] 



Figure 1.1. Stellar Parallax. (1) The Earth orbits the Sun, and a nearby star shifts its position 
with respect to the background stars. The parallax (p) of the star is the angular measurement of 
the Earth's orbit as seen from the star. (2) The closer the star, the greater the parallax angle (PA). 


from the Earth to the Sun, then the star’s distance is 1 parsec (pc) — “the 
distance of an object that has a parallax of one second of arc.” This is the 
origin of the term parsec, which is the unit of distance used most frequently in 
astronomy. 1 

The distance ( d ) of a star in parsecs is given by the reciprocal of its parallax 
( p ), and is usually expressed as thus: 


d = 


1 

P 


Thus, using the above equation, a star with a measured parallax of 0.1 arcseconds 
is at a distance of 10 pc, and another with a parallax of 0.05 arcseconds is 20 pc 
distant. 


Box 1.1: Relationship between Parallax 
and Distance 


d = 


i 

p 


d = the distance to a star measured in pc 
p — the parallax angle of the measured star in arcseconds 

This simple relationship is a significant reason that most astronomical distances are 
expressed in parsecs, rather than light years (l.y.). The brightest star in the night sky is 




Tools of the Trade 


3 


Sirius (a Cams Majoris), which has a parallax of 0.379 arcseconds. Thus, its distance 
from the Earth is: 


, 1 1 

d — = = 2.63 pc 

p 0.379 r 


Note that lpc is equivalent to 3.26 l.y. This distance can also be expressed as: 


d = 2.63 x 


3.26l.y. 
1 pc 


8.61.y. 


Surprisingly, all known stars have a parallax angle smaller than 1 arcsecond, 
and angles smaller than 0.01 arcseconds are very difficult to measure from Earth 
due to the effects of the atmosphere; this limits the distance measured to about 
100 pc (1/0.01). However, the satellite Hipparcos, launched in 1989, was able 
to measure parallax angles to an accuracy of 0.001 arcseconds, which allowed 
distances to be determined to about 1000 pc. 2 

But, this great advance in distance determination is useful only for relatively 
close stars. Most of the stars in the Galaxy are too far for parallax measurements 
to be taken. Another method must be resorted to. 

Many stars actually alter their brightness (these are the variable stars). Several 
of them play an important role in distance determination. Although we shall 
discuss their properties in far greater detail later, it is instructive to mention 
them here. 

Two types of variable stars are particularly useful in determining distances. 
These are the Cepheid variable stars and RR Lyrae variable stars. 3 Both are 
classified as pulsating variables, which are stars that actually change their 
diameter over a period of time. The importance of these stars lies in the 
fact that their average brightnesses, or luminosities, 4 and their periods of 
variability are linked. The longer the time taken for the star to vary its 
brightness (the period), the greater the luminosity. This is the justifiably 
famous Period-Luminosity relationship. 5 The period of a star is relatively easy 
to measure, and this is something that many amateur astronomers still do. 
Once the period has been measured, you can determine the star’s luminosity. 
By comparing the luminosity, which is a measure of the intrinsic brightness of 
the star, with the brightness it appears to have in the sky, its distance can be 
calculated. 6 Using Cepheid as a reference, distances of up to around 60 million 
l.y. have been determined. 

A similar approach is taken with the RR Lyrae stars, which are less luminous 
than Cepheids and have periods of less than a day. These stars allow distances 
to about 2 million l.y. to be determined. 

Another method of distance determination is that of spectroscopic parallax, 
whereby determining a star’s spectral classification can lead to a measure of its 
intrinsic luminosity, which can then be compared with its apparent brightness 
to determine its distance. 



4 


Astrophysics is Easy 


There are other distance determination methods used for the objects farthest 
from us — galaxies. These methods include the Tully Fisher method and the very 
famous Hubble Law. 

All of these methods — Cepheid variable, Tully Fisher, and the Hubble Law — 
will be addressed in greater detail later in the book. 

A final note on distance determination is in order. Do not be fooled into 
thinking that these methods produce exact measurements. They do not. A small 
amount of error is inevitable. This error is usually about 10 or 25%, and even an 
error of 50% is not unheard of. Remember that a 25% error for a star estimated 
to be at a distance of 4000 l.y. means it could be anywhere from 3000 to 5000 l.y. 
away. Table 1.1 lists the 20 nearest stars. 

Let us now discuss some of the nearest stars in the night sky from an obser- 
vational point of view. The list (Table 1.1) is by no means complete but includes 
those stars that are easily seen. Many of the nearest stars are very faint, and thus 
present an observing challenge; they are not included here. 

Throughout the book, I have used the following nomenclature with regard to 
stars: the first item will be its common name, followed by its scientific designation. 
The next item will be its position in right ascension and declination. The final 
item will identify the months when the star is best positioned for observation 
(the month in bold type is the most favorable time of observation). 

The next line will present standard data and information pertinent to the star 
under discussion: its apparent magnitude, followed by its absolute magnitude, 
other specific data relating to the star, and, finally, the constellation in which the 
star resides. 


Table 1.1 

. The 20 nearest stars in the sky 




Star 

Distance, l.y. 

Constellation 

1 

Sun 

.... 

.... 

2 

Proximo Centauri 

4.22 

Centaurus 

3 

Alpha Centauri A 7 

4.39 

Centaurus 

4 

Barnard's Star 

5.94 

Ophiuchus 

5 

Wolf 359 

7.8 

Leo 

6 

Lalande 21185 

8.31 

Ursa Major 

7 

Sirius A 7 

8.60 

Canis Major 

8 

UV Ceti A 7 

8.7 

Cetus 

9 

Ross 1 54 

9.69 

Sagittarius 

10 

Ross 248 

10.3 

Andromeda 

11 

Epsilon Eridani 

10.49 

Eradinus 

12 

HD 217987 

10.73 

Piscis Austrinus 

13 

Ross 1 28 

10.89 

Virgo 

14 

L 789-6 A 7 

11.2 

Aquarius 

15 

61 Cygni A 

11.35 

Cygnus 

16 

Procyon A 7 

1 1.42 

Canis Minoris 

17 

61 Cygni B 

1 1.43 

Cygnus 

18 

HD1 73740 

1 1.47 

Draco 

19 

HD 1 73739 

1 1.64 

Draco 

20 

GX Andromadae 7 

1 1.64 

Andromeda 


Tools of the Trade 


5 


1.1.1 The Nearest Stars to Us 9 

Proxima Centauri V645 Cen 14 h 29.7 m — 62°41' Mar-Apr-May 
11.01 v m 8 15.45M 4.22 l.y. 0.772 " Centaurus 

This is the second-closest star to the Earth and the closest star to the Solar 
System, and thus it is included albeit faint. It is a red dwarf star and also a flare 
star with frequent bursts, having maximum amplitude of around one magnitude. 
Recent data indicate that it is not, as previously thought, physically associated 
with a Centauri, but is, in fact, on a hyperbolic orbit around the star and just 
passing through the system. 

Sirius A[a Canis Majoris] 06 h 45.1 m —16° 43' Dec-Jan-Feb 

-1.44 m 1.45M 8.6 l.y. 0.379" Canis Major 

Sirius, also known as the Dog Star, is a lovely star to observe. It is the sixth-closest 
and brightest star in the sky. It is famous among amateur astronomers for the 
exotic range of colors it exhibits due to the effects of the atmosphere. It also 
has a dwarf star companion — the first to be discovered. A dazzling sight in any 
optical device. 

Procyon a Canis Minoris 07 h 39.3 m +56°13' Dec-Jan-Feb 
0.40 m 2.68M 11.41 l.y. 0.283" Canis Minor 

Procyon is the fifteenth-nearest star and the eighth brightest. Like its neighbor 
Sirius, Procyon has a white dwarf companion star, but it is not visible through 
amateur telescopes. 

Barnard's Star[HD21185] 17 h 57.8 m +4°38' Apr-May-Jun 

9.54 m 13.24M 5.94 l.y. 0.549" Ophiuchus 

The third-closest star is a red dwarf. What makes this star so famous is that it 
has the largest proper motion of any star 10 — 0.4 arcseconds per year. Barnard’s 
Star, also known as Barnard’s Runaway Star, has a velocity of 140 km per second; 
at this rate, it would take 150 years for the star to move the distance equivalent 
to the Moon’s diameter across the sky. It has also been thought that the star 
belonged to the Galaxy’s Halo Population. 

61 Cygni A V1803Cyg 21 h 06.9 m +38°45' Jul-Aug-Sep 
5.20 v m 7.49M 11.35 l.y. 0.287" Cygnus 

This is a very nice double star with separation 30.3 arcseconds and a PA of 150° 
(see section 3.7). Both stars are dwarfs and have a nice orange color. It is famous 
for being the first star to have its distance measured successfully, by F. W. Bessel 
in 1838, using stellar parallax. 

GX And Grb34 00 h 18.2 m +44°01' Aug-Sep-Oct 

8.09 v m 10.33M 11.65 l.y. 0.280" Andromeda 


6 


Astrophysics is Easy 


This is one of a noted red dwarf binary system with the primary star itself a 
spectroscopic double star. Also known as Groombridge 34 A, GX And is located 
about 1 /4° north of 26 Andromedae. 

Lacille HD 217987 11 23 h 05.5 m -35°52' Aug-Sep-Oct 
7.35m 9.76M 10.73 l.y. 0.304" Pisces Austrinus 

This is a red dwarf star with the fourth-fastest proper motion of any known 
star, traversing a distance of nearly 7 arcseconds a year and thus would take 
about 1000 years to cover the angular distance of the full Moon, which is 1/2°. 
Lacille is in the extreme southeast of the constellation, about 1° SSE of 77 Pisces 
Austrinus. 


UV Ceti L726-8A 01 h 38.8 m -17°57' Sep-Oct-Nov 
12.56 v m 15.42M 8.56 l.y. 0.381" Cetus 

The seventh-closest star is a red dwarf system, which is rather difficult, but not 
impossible, to observe. The prefix UV indicates that the two components are 
flare stars; the fainter star is referred to in older texts as “Luytens Flare Star ” 
after its discoverer, W. J. Luyten, who first observed it in 1949. 

Epsilon Eridani HD 22049 03 h 32.9 m -09°77' Oct-Nov-Dec 

3.72m 6.18M 10.49 l.y. 0.311" Eridanus 

The tenth-closest star is a naked-eye object. Recent observations indicate that 
there may be an unseen companion star with a very small mass, approximately 
0.048 that of the Sun. 


1 .2 Brightness and Luminosity 


There is an immense number of stars and galaxies in the sky, and for the most 
part they are powered by the same process that fuels the Sun. This does not mean 
that they are all alike. Stars differ in many respects, such as mass, size, and so 
on. One of the most important characteristics is their luminosity, L. Luminosity 
is usually measured in watts (W), or as a multiple of the Sun’s luminosity, 12 L Q . 
This is the amount of energy that the star emits each second. However, we cannot 
measure a star’s luminosity directly because its brightness as seen from Earth 
depends on its distance as well as its true luminosity. For instance, a Centauri 
A and the Sun have similar luminosities, but, in the night sky, a Centauri A is a 
dim point of light because it is about 280,000 times farther from the Earth than 
the Sun. 

To determine the true luminosity of a star, we need to know its apparent 
brightness, which we define as the amount of light reaching the Earth per unit 



Tools of the Trade 


7 


area . 13 As light moves away from the star, it will spread out over increasingly 
larger regions of space, obeying what is termed an inverse square law. If the sun 
were to be viewed at a distance twice that of the Earth, then it would appear 
fainter by a factor of 2 2 = 4. If we view it from a distance 10 times that of the 
Earth, it would appear 10 2 times fainter. If we were to observe the Sun from the 
same location as a Centauri A, it would be dimmed by 270, 000 2 , which is 70 
billion times! 

The inverse square law describes the amount of energy that enters, say, your 
eye or a detector. Imagine an enormous sphere of radius d, centered on a star. 
The amount of light that will pass through a square meter of the sphere’s surface 
is the total luminosity (L) divided by the total surface area of the sphere. Now, 
as the surface area of a sphere is given by the formula 4 Trd 2 , you will understand 
that, as the sphere increases, d increases, and so does the amount of luminosity. 
You may understand now why the amount of luminosity that arrives at the Earth 
from a star is determined by the star’s distance. 

This quantity, the amount of energy that arrives at our eye, is the apparent 
brightness mentioned earlier (sometimes just called the brightness of a star). It 
is measured in watts per square meter (W/m 2 ). 


Box 1.2: The Luminosity Distance Formula 

The relationship between distance, brightness, and luminosity is given as: 


47id 2 


where b is the brightness of the star in W /m 2 
L is the star’s luminosity in W 
d is the distance to the star in m 

Example: 

Let us apply this formula to Sirius, which is at a distance of 8.6 l.y. 

[Note: 1 l.y. is 9.46 x 10 15 m; thus 8.6 l.y. is 8.6 x 9.46 x 10 15 = 8. 14 x 10 16 m] 

, 3.86 x 10 26 W 

b = 7 

4 it (8.14 x 10 16 m) 2 

b= lx 10“ 7 W/m 2 


This means that, say, a detector of area 1 m 2 (possibly a reflecting telescope) will 
receive approximately one-ten millionth of a watt! 


Astronomers measure a star’s brightness with light-sensitive detectors, and 
this procedure is called photometry. 



8 


Astrophysics is Easy 


Box 1.3: Luminosity, Distance, 
and Brightness 

To determine a star’s luminosity, we need to know its distance and apparent brightness. 
We can achieve this quite easily by using the Sun as a reference. First, let us rearrange 
the formula thus: 

L = 4vd 2 b 

Now, applying this equation to the Sun, whose luminosity is given as L 0 and distance 
is d 0 , which is equal to 1 AU, the Sun’s apparent brightness (b 0 ) is: 

l q = 4 ^ d o 2b G 

Now let us take the ratio of the two formulas: 

(L = 47 Td 2 b)/(L Q = 47td 0 2 b 0 ) 

which gives us: 

L/L 0 = (d/d Q ) 2 b/b Q 

Therefore, all we need to know to determine a star’s distance is how far it is as 
compared with the Earth-Sun distance, given as d/ d 0 , and how bright it is as compared 
with that of the Sun, given as b/b Q . 

Example: 

Let Star 1 be at half the distance of Star 2, and appear twice as bright. Compare the 
luminosities. First, d l /d 2 = 1/2; similarly, b l /b 2 = 2. Then: 



This means that Star 1 has only half the luminosity of Star 2, but it appears brighter 
because it is closer to us. 


1 .3 Magnitudes 


Probably the first thing anyone notices when looking at the night sky is that the 
stars differ in brightness. A handful are bright, a few others are fairly bright, 
and the majority are faint. This characteristic, the brightness of a star, is called 
its magnitude (also refers to any other astronomical object that is observed with 
the naked eye). Magnitude is one of the oldest scientific classifications used 
today, and it was coined by the Greek astronomer Hipparchus. He classified 
the brightest stars as first-magnitude stars; those that were about half as bright 
as first-magnitude stars were called second-magnitude stars, and so on; the 




Tools of the Trade 


9 


sixth-magnitude stars were the faintest he could see. 14 Today, we can see the 
fainter stars, and so the magnitude range is even greater, down to thirtieth- 
magnitude. Because the scale relates to how bright a star appears to an observer 
on Earth, the term is more correctly called apparent magnitude, and is denoted 
by m. 

You may have noticed by now that this is a confusing measurement because the 
brighter objects have smaller values; for example, a star of apparent magnitude 
+4 (fourth-magnitude) is fainter than a star of apparent magnitude +3 (third- 
magnitude). Despite the confusion in its usage, apparent magnitude is used 
universally, and so we are stuck with it. A further point to note is that the 
classification of stars has undergone revision since Hipparchus’s day, and an 
attempt was made to put the scale on a scientific footing. In the 19th century, 
astronomers accurately measured the light from stars, and they were able to 
determine that a first-magnitude star is about 100 times brighter than a sixth- 
magnitude star, as observed from Earth. To put it another way, it would take 100 
sixth-magnitude stars to emit the light of one first-magnitude star. The definition 
for magnitude scale was then stated to be thus: a difference of 5 magnitudes 
corresponds exactly to a factor of 100 in brightness (see Table 1.2). A difference 
in magnitude of 1 corresponds to a factor of 2.512 in brightness. This is shown 
by the following calculation: 

2.512 x 2.512 x 2.512 x 2.512 x 2.512 = (2.512) 5 = 100 

Using this modern scale, several objects now have negative magnitude. Sirius, 
the brightest star in the sky, has a value of —1.44 m; Venus (at brightest) has 
—4.4 m, the full Moon has —12.6 m, and the Sun has —26.7 m. 


Table 1 .2. Magnitude and brightness 
ratio difference 


Magnitude Difference Brightness Ratio 


0.0 

1.0 

0.1 

1.1 

0.2 

1.2 

0.3 

1.3 

0.4 

1.45 

0.5 

1.6 

0.7 

1.9 

1 

2.5 

2 

6.3 

3 

16 

4 

40 

5 

100 

7 

630 

10 

10,000 

15 

1 ,000,000 

20 

10,000,000 


10 


Astrophysics is Easy 


Box 1 .4: Apparent Magnitude 
and Brightness Ratio 

Both apparent magnitude (m) and absolute magnitude (M) are used by astronomers, 
and there are several relationships between them. Consider two stars, s l and s 2 , which 
have apparent magnitudes m 1 and m 2 and brightnesses b l and b 2 , respectively. The 
relationship between them can be written as: 


m l — m 2 = — 2.5log 

This means that the ratio of their apparent brightnesses ( b 1 /b 2 ) corresponds to the 
difference in their apparent magnitudes ( m 1 -m 2 ). 

Example: 

Sirius A has a magnitude of —1.44, while the Sun has a magnitude of —26.8. The 
ratio of their brightnesses is thus: 


m l —m 2 = — 2.5log 


b i 


— 1.44— (—26.8) = — 2.5log 1 &sin “ s 


- 10.21 




b sirius \ 

b S u„ ) 


10~ 101 = 7.9 x 10~ n = 1/1.32 x 10 10 


Thus, Sirius A appears 13,200,000,000 times fainter than the Sun, even though it is 
more luminous (as well as more distant). 


The apparent magnitude scale does not tell us whether a star is bright because it 
is close to us, or faint because it is small or distant; all that the classification tells 
us is the apparent brightness of the star — that is, the star’s brightness as observed 
with the naked eye or through a telescope. A more precise definition may be that 
of the absolute magnitude ( M ) of a star; it is defined as the brightness an object 
would have at a distance of 10 pc. This is an arbitrary distance, derived from 
stellar parallax, the technique mentioned earlier; nevertheless, it quantifies the 
brightness of stars in a more rigorous way. 16 As an example, Deneb, a lovely star 
of the summer sky, in the constellation Cygnus, has an absolute magnitude of 
—8.73, making it one of the intrinsically brightest stars, while Van Biesbroeck’s 
star has a magnitude of +18.6, making it one of the intrinsically faintest stars 
known. Table 1.3 shows 20 brightest stars. 



Tools of the Trade 


11 


Table 1.3. 

The 20 brightest stars 

in the sky 



Star 

Apparent Magnitude, m 

Constellation 

1 

Sirius 

— 1 .44 v 15 

Canis Major 

2 

Canopus 

— 0.62 v 

Carina 

3 

Alpha Centauri 

-0.28 

Centaurus 

4 

Arcturus 

-0.05 v 

Bootes 

5 

Vega 

0.03 v 

Lyra 

6 

Capella 

0.08 v 

Auriga 

7 

Rigel 

0.1 8 V 

Orion 

8 

Procyon 

0.40 

Canis Minor 

9 

Achernar 

0.45 v 

Eradinus 

10 

Betelgeuse 

0.45 v 

Orion 

1 1 

Hadar 

0.61 v 

Centaurus 

12 

Altair 

0.76 v 

Aquila 

13 

Acrux 

0.77 

Crux 

14 

Aldebaran 

0.87 

Taurus 

15 

Spica 

0.98 v 

Virgo 

16 

An tares 

1.05 v 

Scorpius 

17 

Pollux 

1.16 

Gemini 

18 

Fomalhaut 

1.16 

Piscis Austrinus 

19 

Becrux 

1.25 v 

Crux 

20 

Deneb 

1.25 

Cygnus 


Box 1 .5: Relationship Between Apparent 
Magnitude and Absolute Magnitude 

The apparent magnitude and absolute magnitude of a star can be used to determine 
its distance, the formula for which is: 

m — M= 5 log d — 5 

where m = the star’s apparent magnitude 
M = the star’s absolute magnitude 
d = the distance to the star (in pc) 

The term m — M is referred to as the distance modulus. 

Example: 

Sirius is at a distance of 2.63 pc and has an apparent magnitude of — 1 .44. Its absolute 
magnitude can be calculated thus: 

m — M = 5 log d — 5 
M = m — 5 log d + 5 
— 1 .44 — 5 log(2.63) + 5 
M ~ 1.46 



12 


Astrophysics is Easy 


1 .3. 1 The Brightest Stars 

Below is a list of some of the brightest stars in the sky. It is by no means 
complete. For those interested in observing additional bright stars, I recommend 
the accompanying volume to this book. Several of the brightest stars have already 
been mentioned in the section “The Nearest Stars.” For the sake of clarity and 
space, they will not be repeated here. 

Pollux (i Gem 07 h 45.3 m +28°02' Dec-Jan-Feb 

1.16m 1.09M 33.72 l.y. Gemini 

This is the brighter of the two famous stars in Gemini, the other being 
Castor. 


Becrux pCrucis 12 h 47.7 m -59°41' Mar-Apr-May 

1.25 v m -3.92 M 352.1 l.y. Crux 

This star lies in the same field as the glorious Jewel Box star cluster. It is a 
pulsating variable star with a very small change in brightness. 

Spica aVirginis 13 h 25.2 m —11° 10' Mar-Apr-May 

0.98 v m -3.55 M 262 l.y. Virgo 

The fifteenth-brightest star is a large spectroscopic binary with the companion 
star lying very close to it, and thus eclipsing it slightly. Spica is also a pulsating 
variable star, though the variability and pulsations are not visible with amateur 
equipment. 

Hadar pCentauri 14''03.8 m -60°22' Mar-Apr-May 

0.58 v m — 5.45 M 525 l.y. Centaurus 

This is the eleventh-brightest star in the sky, and it is invisible to northern 
observers because of its low latitude (lying only 4.5° from a Centauri). 
It has a luminosity that is an astonishing 10,000 times that of the Sun. 
A white star, it has a companion of magnitude 4.1, but it is a difficult 
double to split, as the companion is only 1.28 arcseconds from the 
primary. 


Arcturus aBootis 14''15.6 m +19°U' Mar-Apr-May 

— 0.16 v m — 0.10M 36.7 l.y. Bootes 

The fourth-brightest star in the sky, Arcturus is the brightest star north of the 
celestial equator. It has a lovely orange color and is notable for its peculiar 
motion through space. Unlike most stars, Arcturus is not travelling in the plane 
of the Milky Way, but instead it is circling the Galactic center in a highly inclined 
orbit. Calculations predict that it will swoop past the Solar System in several 
thousand years, moving towards the constellation Virgo. Some astronomers 
believe that, in as little as half a million years, Arcturus will disappear from 


Tools of the Trade 


13 


naked-eye visibility. At present, it is about 100 times more luminous than 
the Sun. 

Rigil Kentaurus a Centauri 14 h 39.6 m -60°50' Apr-May-Jun 

—0.20 m 4.07 M 4.39 l.y. Centaurus 

The third-brightest star in the sky, Rigil Kentaurus is in fact a part of a triple 
system, with the two brightest components contributing most of the light. The 
system contains the closest star to the Sun, Proxima Centauri. The group also has 
a very large proper motion (its apparent motion in relation to the background). 

Antares a Scorpii \6 h 29A m — 26°26' Apr-May-Jun 

1.06 v m — 5.28M 604 l.y. Scorpius 

This is a red giant star with a luminosity 6000 times that of the Sun and a diameter 
hundreds of times larger than the Sun’s. What makes this star especially worth 
watching is the vivid color contrast that is observed between it and its companion 
star. The star is often described as vivid green when seen with the red of Antares. 
The companion has a magnitude of 5.4, PA of 273°, and lies 2.6" away. 

Vega aLyrae 18 ,, 36.9 m +38°47' Jun-Jul-Aug 

0.03 v m 0.58 M 25.3 l.y. Lyra 

This is the fifth-brightest star, familiar to northern observers, located high in the 
summer sky. Although similar to Sirius in composition and size, Vega is three 
times as distant, and thus appears fainter. Often described as having a steely-blue 
color, it was one of the first stars observed to have a disc of dust surrounding 
it — a possible proto-solar system in formation. Vega was the Pole Star some 
12,000 years ago, and will be again in another 12,000 years. 

Altair cxAquilae 19 ,! 50.8 m +08°52' Jun-Jul-Aug 

0.76 v m 2.20 M 16.77 l.y. Aquila 

The twelfth-brightest star, Altair has the honour of being the fastest-spinning 
of the bright stars, completing one revolution in approximately 6 1/2 hours. 
Such a high speed deforms the star into what is called a flattened ellipsoid, 
and it is believed that, because of this amazing property, the star may have 
an equatorial diameter twice that of its polar diameter. The star’s color has 
been reported to be completely white, although some observers see a hint of 
yellow. 

Fomalhaut a Pisces Austrini 22 h 57.6 m — 29°37' Aug-Sep-Oct 

1.17 m 1.74 M 25.07 l.y. Pisces Austrinus 

The eighteenth-brightest star is a white one, which often appears reddish to 
northern observers due to the effect of the atmosphere. It lies in a barren area 
of the sky, and it is remarkable that a star close to it, which is not bound 
gravitationally yet lies at the same distance from the Earth, is moving through 
space in a manner and direction similar to Fomalhaut. It has been suggested 


14 


Astrophysics is Easy 


that the two stars are remnants of a star cluster or star association that has 
long since dispersed. This orange star of magnitude 6.5 lies about 2° south of 
Fomalhaut. 

Achernar a Eridani 0l''37.7"' -57° 14' Sep-Oct-Nov 

0.45 v m —2.77 M 144 l.y. Eridanus 

The ninth-brightest star in the sky lies at the southernmost end of the constel- 
lation and too far for northern observers. Among the brightest stars, it is one of 
the very few with the designation “p” in its stellar classification, indicating that 
it is a “peculiar” star. 

Aldebaran a Tauri 04 ,! 35.9 m +16°31' Oct-Nov-Dec 

0.87 m — 0.63 M 65.11 l.y. Taurus 

The fourteenth-brightest star appears to be located in the star cluster Hyades; 
however, it is not physically in the cluster at all, lying twice as close as the cluster 
members. This pale-orange star is approximately 120 times more luminous than 
the Sun. It is a double star, too, but very difficult to separate due to the extreme 
faintness of the companion. The companion star, a red dwarf star of magnitude 
13.4, lies at a PA of 34° and at a distance of 121.7". 

Rigel P Orionis 05 ,! 14.5 m — 08°12' Nov-Dec-Jan 

— 0.18 v m -6.69 M 773 l.y. Orion 

The seventh-brightest star in the sky, Rigel is in fact brighter than a Orionis. 
This supergiant star is one of the most luminous stars in our part of the Galaxy, 
almost 560,000 times more luminous than the Sun but at a greater distance than 
any other nearby bright star. Often described as a bluish star, it has a tremendous 
mass — about 50 times that of the Sun and about 50 times its diameter. It has 
a close bluish companion at a PA of 202°, apparent magnitude 6.8, and at a 
distance of 9 arcseconds, which should be visible with a 15 cm telescope, or even 
a smaller one under excellent observing conditions. 

Capella aAurigae 05 h 16.7 m +46°00' Nov-Dec-Jan 

0.08 v m — 0.48 M 42 l.y. Auriga 

The sixth-brightest star in the sky is, in fact, a spectroscopic double, although 
it cannot be split in a telescope; however, it has a fainter lOth-magnitude star 
about 12 arcseconds to the south-east at a PA of 137°. This is a red dwarf star, 
which is itself a double (only visible in larger telescopes). So, Capella is in fact a 
quadruple system. 

Betelgeuse a Orionis 05 h 55.2 m +07°24' Nov-Dec-Jan 

0.45 v m -5.14 M 427 l.y. Orion 

The tenth-brightest star in the sky, and a favorite to most observers, this orange- 
red star is a giant variable with an irregular period. Recent observations by the 
Hubble Space Telescope have shown that it has features on its surface that are 
similar to Sunspots but much larger, covering perhaps a tenth of the surface. 


Tools of the Trade 


15 


It also has a companion star, which may be responsible for the non-spherical 
shape it exhibits. Although a giant star, it has a very low density and a mass 
20 times that of the Sun, which together mean that its density is in fact just 
0.000000005 times that of the Sun. 


1 .4 Color 


When we look up into the sky at night, we see many stars; all of them are generally 
white. There are, of course, a few that exhibit distinct colors — Betelgeuse (a Orionis) 
is most definitely red, as is Antares (a Scorpi); Capella (a Aurigae) is yellow; 
and Vega ( a Lyrae) is steely blue. However, for the most part, there does not 
seem to be any great variation in color. Look through binoculars or a telescope, 
and the situation changes dramatically. 17 Variations in color and hue abound! 18 

The color of a star is determined by its surface temperature. A red star has a 
lower temperature than that of a yellow star, which in turn has a lower temper- 
ature than that of a blue star. This is an example of what is called the Wien 
Law (See Box 1.6). The law states that low-temperature stars emit most of their 
energy in the red to infrared part of the spectrum, while much hotter stars 
emit in the blue to ultraviolet part of the spectrum. Some very hot stars emit 
most of their energy in the ultraviolet, and so in fact we see only a fraction of 
their light. Furthermore, many stars emit nearly all of their light in the infrared, 
and so we do not see them at all. Surprisingly, these low-mass (to be discussed 
later), low-temperature stars make up about 70% of the stars in our galaxy, 
but you would never ever see them by going out and observing the sky on a 
clear night. 

An important point to notice here is that hotter objects emit more energy 
at all wavelengths due to the higher average energy of all the photons. This 
is illustrated in Figure 1.2. The graphs demonstrate how the light from three 
different stars is distributed, depending on the stars’ temperature. The colored 
block represents the visible part of the spectrum. The first plot shows the light 
that would be measured from a colored star of about 3000 K. Note that the curved 
line peaks at about llOOnm, which would make the star appear red. The second 
plot shows a star of about 5500 K (similar to the Sun’s temperature), which peaks 
in the middle of the visible spectrum, thus looking yellowish. The final plot 
illustrates a very hot star, of 25,000 K; it peaks at about 400 nm, and so it will 
appear blue. Thus, the color of a star, from an astronomical viewpoint, depends 
on where the peak of its curve lies; a short wavelength (the left part of the plot) 
indicates a hot, bluish-white star, while a longer wavelength (the right part of 
the plot) indicates a cool, reddish-orange star. The Sun peaks at the green part 
of the spectrum; since there is a mixture of light from all the other parts of the 
visible spectrum — the blues, reds, and yellows — we actually observe the Sun as 
being yellowish-white. 

An interesting thing to observe is that a few stars are so hot, possibly in 
millions of degrees, that they emit energy at very short wavelengths. In fact, they 
radiate X-rays. These are neutron stars! 



16 


Astrophysics is Easy 


Intensity 



Intensity 




Wavelength (nm) 


Figure 1.2. Relationship between color and temperature. 




Tools of the Trade 


17 


Note that when we speak of a star’s temperature, we are referring to its surface 
temperature. The internal temperature cannot be measured directly, and it is 
usually determined from theoretical temperatures. So, when you read that a star’s 
temperature is 25,000 K, it refers to the surface temperature. 19 


Box 1 .6: The Wien Law 


Wien’s Law can be stated as: 


2, 900, 000 

= nm 

max T(Kelvin) 

Example: 

Two stars, a Canis Majoris and o Ceti, have a temperature of 9200 K and 1900 K, 
respectively. What are their peak wavelengths? 


N 2,900,000 . , , . , 

= nm = 315 nm (i.e., in the ultraviolet ) 

max 9200(Kelvin) 


and 


= nm = 1526 nm (i.e., in the infrared 21 ) 

1900(Kelvin) 


Sirius emits a lot of light in the ultraviolet, even though it shines brightly white. 


Knowing a star’s temperature helps to determine many other characteristics. 
One scientific description of a star’s color is based on the stellar classification, 
which in turn is dependent upon the star’s chemical composition and temper- 
ature. A term commonly used by astronomers is color index. It is determined 
by observing a star through two filters, the B and V filters, which correspond to 
the wavelengths 440 and 550 nm, respectively, and by measuring its brightness. 
Subtracting the two values obtained, B and V, produces the color index. Usually, 
a blue star will have a negative color index (e.g., —0.3); orange-red stars will have 
a value greater than 0.0 and upwards to about 3.00, and even greater for very red 
stars (M6 and greater) [see section 1.8]. 

Having discussed the colors of stars, let us now look at some examples. I have 
chosen a representative selection of bright stars. There are, of course, literally 
thousands of other visible colored stars. The section “The Brightest Stars” offers 
many examples of stars that exhibit distinct colors. In addition, many double 
stars (not mentioned here) show very distinct hues and tints. The nomenclature 
used here is the same as that used previously, except for the addition of the star’s 
temperature and color. 22 



18 


Astrophysics is Easy 


1 .4. 1 Colored Stars 

Bellatrix yOri 05 h 26.2 m +06°21' Nov-Dec-Jan 

1.64 m -2.72 M 21, 450 K Blue Orion 

Also known as the Amazon Star, Bellatrix appears steely blue. Some observers 
report a faint nebulosity associated with the star, but it may just be a part of 
general nebulosity that envelopes much of Orion. 

Merope 23Tau 03 ,, 46.3 m +23°57' Oct-Nov-Dec 

4.14 m — 1.07 M 10, 600 K Blue Taurus 

Located within the Pleiades star cluster, it gives a breathtaking and spectacular 
view when seen through binoculars, and the cluster is a highlight of the night sky. 
Almost all the stars in this cluster are worth observing for their lovely steely-blue 
color. 


Regulus a Leo 10 ,! 08.3 m +11°58' Jan-Feb-Mar 

1.36 m — 0.52 M 12, 000 K Blue-white Leo 

a Leonis is the handle of the Lion’s sickle. It is an easy double star; the companion, 
an 8th magnitude, orange-red color, is about 3' away. 

Acrux aCrucis I2 h 26.6 m — 63°06' Feb-Mar-Apr 

0.72 10 m -4.19 M 28, 000/26, 000 K White Crux 

Acrux is a double star with components 4" apart. Both stars are almost of the 
same magnitude: 1.4 for ct 1 and 1.9 for a 2 . The colors of the stars are white and 
bluish-white, respectively. 


Zubeneschamali PLib 15 A 17.0 m —09° 23' Apr-May-Jun 

2.61m — 0.84 M 11, 000 K Green! Libra 

This is a mysterious star for two reasons: historical records state that it was much 
brighter than it appears today, and observers of the past 100 years have declared 
that it is greenish or pale emerald in color. It is one of the rare green-colored 
stars! 


The Sun Jan-Dec 

—26.78 m 4.82 M 5800 K Yellow The Zodiac 

The Sun is our closest star, without which no life would have evolved on Earth. It 
is visible every day throughout the year, unless you happen to live in the UK. DO 

NOT OBSERVE THE SUN THROUGH ANY KIND OF OPTICAL EQUIPMENT. 


Tools of the Trade 


19 


Garnet Star (jl Cep21 h 43.5 m +58°47' Jul-Aug-Sep 

4.08 v m — 7.3M 3500 K Orange Cepheus 

Located on the north-eastern edge of the nebulosity IC1396, Garnet Star, named 
by William Herschel, is one of the reddest stars in the sky, having a deep 
orange or red color, seen against a backdrop of faint white stars. It is a 
pulsating red giant star, with a period of about 730 days and brightness from 
3.4 to 5.1 m. 

Hind's Crimson Star R Leporis 04' , 59.6" 1 -14° 48' Nov-Dec-Jan 
7.71 v m 1.08M 3000 K 23 Red Lepus 

This star, a classic long-period variable, has a period of about 432 days, and 
varies in brightness between 6.0 and 9.7 m. At maximum brightness, it displays 
the famous ruddy color that gave it its name. It was discovered in 1845 by J. R. 
Hind, who described its color as “intense smoky red.” This may be the reddest 
star in the sky. It is also an AGB star (see section 3.15). 


1 .5 Size and Mass 


Stars are at an immense distance from the Earth; no matter how much we magnify 
a star’s image, it will, in all but a handful of cases, 24 remain just a point of 
light. So how do we determine the size of a star? The answer is quite simple: by 
measuring both the luminosity (derived from its distance and brightness) and 
the surface temperature (determined from its spectral type); it is just a matter of 
manipulating numbers with a few formulas. Using this technique, astronomers 
have discovered that many stars are much smaller than the Sun, while many 
others are thousands of times larger. 

To accurately determine a star’s size, a physical law called the Stefan- 
Boltzmann Law is used. We will not bother looking at how this law came about, 
but simply quote it and show how it is used (see Box 1.7). The law states that 
the amount of energy that a star radiates per second from a square meter of 
its surface 25 is proportional to the fourth power of the temperature (T) of its 
surface. Do not let the complexity of this statement distract you. It just tells us 
that the energy flux (F) is proportional to the temperature, which may make 
sense to you when you think it over. A cool object has lower thermal energy than 
a hot object. 

Now recall what we discussed earlier, that the luminosity of a star is a measure 
of the energy emitted from its surface every second. This luminosity is, in fact, 
the flux (F) multiplied by the number of square meters there are on the star’s 
surface. If we assume that most stars are spherical (which is not as silly as it 
sounds because a few stars are not spherical!), then the quantity highlighted 
in the previous sentence is in fact the surface area of the star. This is given 
by a very simple formula, which most of us already know: 4ttR 2 , where R is 
the radius of the star (taken as the distance from the center of the star to its 
surface 26 ). 




20 


Astrophysics is Easy 


Box 1 .7: Flux, Luminosity, and Radius 

The flux from a star is given by the Stefan-Boltzmann Law: 

F = aT 4 

The relationship between flux (F), luminosity (L), and radius (R) of a star is: 

L = 47rF 2 oT 4 

where L is the star’s luminosity in watts (W) 

R is its radius in meters (m) 

c r is the Stefan-Boltzmann constant; 5.67 x 10~ 8 W m~ 2 K~ 4 
T is the star’s temperature in Kelvin (K) 


The above equations tell us that a coolish star (one with a low surface temper- 
ature) will have a low flux, but it might be quite luminous because it could have a 
very large radius, and thus a large surface area. In a similar vein, a hot star (one 
with a high surface temperature) might have a low luminosity if it has a small 
radius, which would mean a low surface area. Now you can see that knowing the 
temperature alone does not indicate how luminous a star will be — you need its 
radius, too! 

Although we can now determine such parameters as the radius, temperature, 
luminosity, and brightness of a star, it is often more useful to relate these values 
to that of the Sun. It would be easier for someone to understand if we said that 
a star is about 10 times hotter than Sun rather than saying it is 54,000 K. The 
same applies to L and R. 


Box 1 .8: Even More about Flux, 

Luminosity, and Radius 

We regard the Sun as a typical star, and so we can relate a star’s characteristics to that 
of the Sun. For instance: 

L Q = 4 tt R 0 2 aT 0 4 

where L e is the luminosity of the Sun 
R Q is the Sun’s radius 
T Q is the Sun’s temperature 




Tools of the Trade 


21 


If we divide the luminosity equation for a star by that for the Sun, we get: 

L/ L 0 = (R/R 0 ) 2 (T/T 0 ) 2 

The constants cr and 4tt have now gone, and we can rearrange the formula to read: 

R/R e = (T g /W/L g ) 1/2 

where the factor 1/2 indicates a square root. 

Now, R/Rg is the ratio of the star’s radius to that of the Sun 
T 0 /T is the ratio of the Sun’s temperature to that of the star 
L/L 0 is the ratio of the star’s luminosity to that of the Sun 

Example: 

Sirius has a temperature of about 9200 K and luminosity of about 23 L 0 . 

To determine its ratio: 

/ 5800 \ 2 y— 

R/R 0 = x V23 ~ 2 

' G V 9200 / 

Thus, its radius is about twice that of the Sun. 


1 .5.1 The Biggest Stars 

Let us now discuss some examples of giant stars, particularly those that can be 
seen with the naked eye. 

aHerculis ADS 10418 17'’14.6 m +14°23' May-Jun-Jul 

3.5 V , 5.4 v m — 1.9M Radius: 2.0 AU Hercules 

A lovely color-contrast double (orange and bluish-green), the star lies at a 
distance of about 400 l.y. and is a semi-regular, supergiant variable star. The 
primary star is itself variable, while the secondary is an unresolvable double. 

tit 1 Aurigae HD 44537 06 ,! 24.9 m +49°17' Nov-Dec-Jan 

4.92 v m — 5.43 M Radius: 3.0 (?) AU Auriga 

This star has an incredible luminosity of over 11, 000 L Q . It is an irregular variable 
star, the diameter of which is not known. The star is believed to be about 
4300 l.y. distant. 

TjPersei ADS 2157 02'’50.7 m +55°54' Oct-Nov-Dec 

3.8, 8.5 m —4 M Radius : 2.0 AU Perseus 

Lying at a distance of about 1300 l.y., this is a lovely double star — gold primary 
and blue secondary. The primary is a supergiant with a luminosity of over 
4000 L q . 

VVCephei HD 208816 2\ h 56.6 m +63° 37' Sep-Oct-Nov 

5.11m — 6.93 M Radius: 8.8 AU Cepheus 



22 


Astrophysics is Easy 


This star has a luminosity between 275,000 and 575, 000 L Q , and it lies at a 
distance of 2000 l.y. It is one of the famous eclipsing binary-type variable stars, 
with a period of just 20 years or over. The system consists of an O-type dwarf 
and an M-type supergiant; if placed at the center of the solar system, this giant 
star would extend to the orbit of Saturn! 

KQ Puppis HD 60414 07 h 33.8 m -14°31' Sep-Oct-Nov 

4.82m — 5.25 M Radius: 8.8 AU Puppis 

This star has a luminosity of over 9870 L Q , and it lies at a distance of 3361 l.y. 27 
It is believed to be an irregular variable star. 


1 .6 Star Constituents 


Although we shall cover this topic in far greater detail later in the book, it is 
important that we briefly look at what stars are made of. 

A star is an enormous sphere of hot gases. It is as simple, or as complex, as 
that, whichever way you wish to look at it. Of course, the processes involved in 
making and maintaining a star are, as expected, very, very complex! 

Gases that compose a star are primarily hydrogen (H, the most common 
element in the universe), helium (He), and some other elements. 28 By and large, 
most stars are nearly entirely made of hydrogen, less helium, and very small 
amounts of everything else. This composition is usually about 75% hydrogen, 
24% helium, and the remainder metals. This ratio may change, however, since 
very old stars are nearly all hydrogen and helium with tiny amounts of metals, 
and very new stars can contain as much as 2-3% metals. 

The energy needed to create and maintain a star is produced within the 
star by nuclear fusion. Two immense forces — very high temperature and strong 
gravitational force — are at work. Due to very large mass and concomitantly strong 
gravitational fields, the conditions at the center of the ball of gases are such that 
the temperature may be about 10 million K. At such extremes of pressure and 
heat, nuclear fusion can occur, by which hydrogen is converted into helium. The 
outcome of this nuclear reaction is a tiny amount of energy in the form of gamma 
rays. It may not seem like much, but when billions of these reactions take place 
every second, the amount of energy liberated is quite substantial... enough, in 
fact, to make a star shine! 

As a star ages, it uses up more and more hydrogen in order to keep the 
nuclear reactions going. A by-product of this reaction is helium. After quite some 
time, the amount of hydrogen decreases and helium increases. If conditions are 
right (which include a higher temperature and a large mass), then helium itself 
will start to undergo nuclear fusion at the star’s core. After a very-long time, 
helium, in turn, will produce the element carbon as a by-product of the reaction; 
similarly, if conditions are suitable, carbon, too, will initiate nuclear fusion and 
produce more energy. An important point to emphasize is that each step requires 
a higher temperature to begin nuclear reactions, and if a star does not have 
the conditions necessary to produce this temperature, further reactions will not 



Tools of the Trade 


23 


occur. So, you can see that the “burning” of hydrogen and helium is the source 
of power for nearly all the stars that we see, and the mass of a star determines 
how the reactions will proceed. 


1 .7 Spectra and Spectroscopy 


Let us now look at a tool that is central to the topic of astrophysics — spectroscopy 
and spectra. This is an amazing topic; from just looking at the light from an 
object, we can say how hot it is, how far away it is, in which direction is it 
moving , 29 if it is rotating, and (from all these data) infer its age, mass, how long 
it will live, and so on. In fact, so important is this topic that, from this point on 
in the book, a star will be referred to by its spectral classification. 

Determining a star’s classification is a theoretically easy task, although it may 
be difficult in practice. What is needed is a spectroscope. This is an instrument 
that helps one look at the light from a star in a special way by utilizing either 
a prism or a diffraction grating for analysis. You are probably aware that white 
light is in fact a mixture of many different colors, or wavelengths, and so it is safe 
to assume that the light from a star is a mixture of colors. Indeed it is, but usually 
with an added component. Using a spectroscope mounted at the eyepiece end of 
the telescope , 30 light from a star can be collected and photographed (these days 
with a CCD camera). The result is something called a spectrum. Many amateur 
astronomers are now making some good observations of stars’ spectra. 

Basically, a spectrum is a map of light coming from a star. It consists of all of 
the emitted light, spread out according to wavelength (color), so that different 
amounts of light at different wavelengths can be measured. Red stars have a lot 
of light at the red end of the spectrum, and blue stars have a correspondingly 
larger amount at the blue end. However, an important point to note is that, in 
addition to this light, there will be a series of dark lines superimposed upon 
this rainbow-like array of colors. These are called absorption lines, which are 
formed in the atmosphere of the star. In rare cases, there are bright lines, too, 
called emission lines. Although comparatively rare in stars, these lines are very 
prominent in nebulae. 

The electrons in the atoms located on the surface layers of a star can only 
have very specific energies, just like the specific heights of the rungs of a ladder. 
Sometimes, an electron in an atom of, say, hydrogen can be “knocked” from a 
lower energy level to a higher energy level, maybe by a collision with another 
atom. Eventually, it will fall down to the lower level. The energy that the atom 
loses when the electron returns to its original level must go somewhere, and 
it often goes to emitting a photon of light. This emitted photon has a unique 
property — it has the exact amount of energy that the electron loses, which in 
turn means that the photon has a very specific wavelength and frequency. 

When hydrogen gas is heated to a high temperature, the number of collisions 
between atoms can continually bump electrons to higher energy levels, and an 
emission line spectrum results. This consists of the photons that are emitted as 
each electron falls back to lower levels. 

The origins of the absorption lines are due to the differing amounts of elements 
in the cooler atmosphere of the stars (recall that in addition to hydrogen and 




24 


Astrophysics is Easy 


helium, there are other elements, or metals, present, but in minute quantities). 
Not only are photons emitted, but they may also be absorbed. This process causes 
the electrons to jump up in energy to a higher level. But, this can only happen 
if the photon has the precise amount of energy required. Too much too little 
energy, even a minuscule amount, can cause the photon to not interact with the 
electron. 

In hydrogen gas, an electron moving from level 2 to level 1 will emit a 
photon that has a wavelength of 121.6 nm; an electron absorbing a photon of this 
wavelength will jump from level 1 to level 2. Such jumps from different levels are 
called transitions. Thus, in the above example, an electron undergoes a transition 
from level 1 to level 2, with an absorption of a photon of wavelength 121.6 nm. 
Figure 1.3 shows the allowed energy levels of hydrogen and wavelengths that 
occur for downward transitions. Also shown are the absorption and emission 
spectra. 

Note in Figure 1.3 that the dark absorption lines and bright emission lines 
occur at exactly the same wavelengths, regardless of whether the hydrogen is 


(a) 


ionization 



Level 5 
Level 4 
Level 3 

Level 2 


Level 1 



410.1 434.0 486.1 656.3 


nm nm nm nm 

(b) 



410.1 434.0 486.1 656.3 


nm nm nm 

(c) 


nm 


Figure 1 .3. Allowed energy levels of hydrogen, (a) The wavelengths of various energy level 
transitions in hydrogen. 31 (b) Visible emission line spectra showing transitions that occur from 
high energy levels downward to level 2 for hydrogen, (c) Absorption line spectra showing 
transitions that arise from energy level 2 to higher levels. These absorption and emission lines of 
hydrogen are called the Balmer Lines. 




Tools of the Trade 


25 


emitting or absorbing light. Emission lines are simply the result of downward 
jumps, or transitions, of electrons between the energy levels, while absorption 
lines are upward transitions. 

The energy levels of electrons in each chemical element are unique — a “finger- 
print” that results in each element’s having its own distinct spectral lines. Hydrogen 
is a very simple element, with only 1 electron, but in those elements with many 
electrons and energy levels, the corresponding spectra can be very complex. 

The factor that determines whether an absorption line will arise is the temper- 
ature of a star’s atmosphere. A hot star will have different absorption lines than a 
cool star. The classification of a star is determined by examining its spectrum and 
measuring various aspects of the absorption lines. A very important point that I 
would emphasize is that the observational classification of a star is determined 
primarily by the temperature of the atmosphere and not the core temperature. 
The structure of the absorption lines can be examined, and this gives further 
information on pressure, rotation, and even the presence of a companion star. 


1 .8 Stellar Glassification 


We saw earlier that stars are distinguished by their spectra (and thus temper- 
ature). Let us now think about the spectral type. For historical reasons, a star’s 
classification is designated by a capital letter which is in order of decreasing 
temperature: 32 

OBAFGKMLRNS 

The sequence goes from hot blue stars (types 0 and A) to cool red stars (types 
K, M, and L). In addition, there are rare hot stars called Wolf-Rayet stars (classes 
WC and WN), exploding stars (Q), and peculiar stars ( p ). The star types R, N, 
and S actually overlap the class M, and so R and N have been reclassified as 
C-type stars, the C standing for carbon stars. A new class (L) has recently been 
introduced. 33 Furthermore, the spectral types are divided into ten spectral classes 
beginning with 0, 1,2, 3, and so on up to 9. A class A1 star is hotter than a class 
A8 star, which in turn is hotter than a class F0 star. Furthermore, prefixes and 
suffixes can be used to illustrate additional features: 


a star with emission lines e 

(also called f in some O-type stars) 
metallic lines m 

peculiar spectrum p 

variable spectrum v 


a star with a blue or red shift in the line q 
(e.g., P Cygni stars) 


For more historical reasons, the spectra of the hotter star of types 0, A, and B 
are sometimes referred to as early-type stars, while the cooler ones ( K , M, L, C, 
and S) are later-type. F and G stars are designated intermediate-type stars. 




26 


Astrophysics is Easy 


Because the spectral type is so important, it is instructive to explain further 
how the appearance of a spectrum is affected by its surface temperature. We will 
consider the Balmer lines of hydrogen, mainly because these are by far the easiest 
to understand. Hydrogen gas makes up 75% of a star, yet the Balmer lines do not 
always appear in a star’s spectrum. The Balmer absorption lines are produced when 
an electron undergoes a transition from the 2nd energy level to a higher level by 
absorbing a photon with the correct amount of energy. If, however, the star is 
hotter than 10,000 K, the photons coming from the star’s interior have such a high 
energy that they can easily knock electrons out of hydrogen atoms in the star’s 
atmosphere. This process is called ionization. Now that the hydrogen atom has 
lost its electrons, it cannot produce absorption lines. So, the Balmer lines will be 
relatively weak in the spectra of such hot stars (e.g., type-0 stars, up to type 152). 

On the other hand, if the atmosphere of a star is cooler than 10,000 K, most 
of the hydrogen atoms are in the 1st energy state. Many of the photons passing 
through the atmosphere do not have enough energy to boost the electron from 
the 1st to the 2nd energy level. Therefore, very few atoms will have electrons in 
the 2nd level, and only these electrons will absorb the photons characteristic of 
the Balmer lines. This results in the lines’ being almost absent from the spectrum 
of cool stars, such as M0 and M2 stars. 

For the Balmer lines to be prominent, a star must be hot enough to excite the 
electrons out of level 1 (also known as the ground state), but not so hot that the 
hydrogen becomes ionized. If a star has a surface temperature of around 9,000 K, 
it will have the strongest hydrogen lines (e.g., the A0 to A5 stars). 

The Balmer lines of hydrogen become increasingly prominent as you go from 
type B0 to A0. From A0 through F and G class, the lines weaken and almost fade 
away. The Sun, a G2 star, has a spectrum dominated by lines of calcium and iron. 

Finally, a star can also be classified by its luminosity, which is related to its 
intrinsic brightness, with the following system: 

Supergiants 34 I 
Bright giants II 


Giants III 

Subgiants IV 

Dwarfs V 

Subdwarfs VI 


White dwarfs VII 

It is evident that astronomers use a complex and very confusing system! In 
fact, several classes of spectral type are no longer in use, and the luminosity 
classification is also open to confusion. It will not surprise you to know that 
there is even disagreement among astronomers as to whether, for example, a star 
labelled F 9 should be reclassified as GO! Nevertheless, it is the system generally 
used, and so it will be adhered to here. Examples of classification are: 


Tools of the Trade 


27 


Bootes (Arcturus) K2IIIp 

(3 Orionis (Rigel) B8Ia 

a Aurigae (Capella) G8 III 

P Cygni Bllapeq 

Sun G2V 


I conclude my discussion on spectral classification by explaining what the spectral 
type actually refers to. 35 You may recall that spectral classification was based on 
the detection of absorption lines, which in turn depends on the temperature of 
a star’s atmosphere. Thus, the classification relies on the detection of certain 
elements in a star, giving rise to a temperature determination for that star. The 
classification is summarized in Table 1.4. 

It is interesting to note that the distribution of stars throughout the Galaxy 
may not be what you assume. A casual glance at the stars in the night sky will 
give you several 0- and B-type, a few A-type, some F- and G-type, a smattering 
of K, and more M-types. You may think that this is a fair picture of the type 
of distribution throughout the remainder of the Galaxy. You could be wrong! 
As we shall see in later sections, a vast majority of stars in our Galaxy — over 
72% of them — are faint, cool, and red M - type stars. The bright and hot O-type 
stars are less than 0.005%. For every O-type star, there are about 1.7 million 
M-types! 

Let us now look at a few examples: 

I . 8. 1 The Spectral Sequence 

HD 93 129 A 10''43.9 m -59°33' Jan-Feb-Mar 

7.0 m — 7.0 M 03If Carina 

This is an extraordinary star! This supergiant star, lying at a distance of about 

I I, 000 l.y. shines about 5 million times as brightly as the Sun. With a mass of 
120 M Q , it is believed to be one of the most luminous stars in the entire Galaxy. 

0 Orionis C 0Ori 05 h 35.3 m — 05°23' Nov-Dec-Jan 

4.96 m — 5.04 M 06 Orion 

A member of the famous Trapezium multiple star system in the Orion Nebula, 
this is a fairly new star, maybe several thousand years old, and as a consequence 
most of its light is emitted at ultraviolet wavelengths. It has a temperature of 
about 45,000 K and a diameter 10 times that of the Sun. 

15Monocerotis HD 47839 06 ,! 40.9 m +09°54' Nov-Dec-Jan 

4.66 v m — 2.3 M 07 Monoceros 

Both a visual binary and a variable star, 15 Monocerotis is located in the star 
cluster NGC 2264, which in turn is encased in a diffuse nebula. 

Plaskett'sStar HD 47129 06 ,! 37.4 m +06°08' Nov-Dec-Jan 

6.05 m — 3.54M 08 Monoceros 


Table 1 .4. Spectral Classification 


Q_ 

£ 

O 

X 


O 


u 


_9- h! 

< < < 


! I 

2 § 3 
£ u - 


0 


05 


05 

"t/i E 
0 -iii 

05 O 


3 E 



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c 

0 

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E 


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o 

Os 

00 


00 


CO 


c 0 
O 1 - 

''Cl' 

c 

„ 

Os 

■> 

CO -XC 



‘O 

£ 

00 


o 

o 

u 

l\ 

Os 

V 

violet 

CN 

1 

K 

Os 

o 

”5 

1 0 

§1 

CN 

1 

o 

Os 

CO 

(blue] 

i 

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_o 

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1 

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(red) 

co S 
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A C3. 


o 

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100 

lore 

han 


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Q_ 

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*55 £ 
2 -2 


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K 

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+ 






co' 

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o v 

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£ >v .E 


.c 

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0 

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0 

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~0 


U 


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0 

c 

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c 

0 

u 

Q_ 

“O 



05 

“O 


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p 

O 

O — 

0 

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N 


0 

^ o 

N 

_Q 

< 

ioni 

D 

0 

C 

Q_ 

Q_ 

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"2 0 
E 

ioni 


= b 

o 0 

u > 

£ c" 


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"8 o -s 3 

_h l_ ^3 0 . 


Z> ~Q ^ 

0 >v O 

C _C _Q 


0 Q_ 
< 8 -^ 


o 


o 


Tools of the Trade 


29 


Plaskett’s Star is actually composed of two stars and is a spectroscopic binary 
system with an estimated mass of approximately 110 times that of the Sun, 
making it one of the most massive known objects. 

GammaK Cassiopeiae yCas 00 , '56.7 m +60° 43' Sep-Oct-Nov 

2.15 v m — 4.22M BO IV Cassiopeia 

This peculiar star has bright emission lines in its spectrum, indicating that it ejects 
material in periodic outbursts. It is the middle star of the familiar W-shape of 
Cassiopeia. 


Mirzim (3 CMa 06 A 22.7 m -17°57' Nov-Dec-Jan 

1.98 v m — 3.96 M Bill CanisMajor 

Mirzim is the prototype of a class of variable stars now classified as /3 Cepheid 
stars, which are pulsating variables. The magnitude variation, however, is too 
small to be observed visually. 

Algenib yPeg 00 ,! 13.2 m +15°11' Aug-Sep-Oct 

2.83 v m — 2.22 M B2V Pegasus 

Algenib is a member of the type fiCMa (Canis Majoris) variable star and is the 
south-eastern corner star of the famed square of Pegasus. 

Achernar aEri 01 h 37.7 m —57° 14' Sep-Oct-Nov 

0.45 v m — 2.77 M B3V Eradinus 

Achernar is a hot and blue star. It lies so far south, it can never be seen from the 
UK. 

Aludra r|CMa 07 ,! 24.1 m — 29°18' Dec-Jan-Feb 

2.45m 7.51 M B5 1 CanisMajor 


This is a highly luminous supergiant with an estimated luminosity 50,000 times 
that of the Sun. 

Electra 17Tau 03' ! 44.9"' +24°07' Oct-Nov-Dec 

3.72 m — 1.56 M B6 III Taurus 


Electra is located within the Pleiades star cluster. 

Alcyone -qTauri 03 ,! 47.5 m +24°06' Oct-Nov-Dec 

2.85 m — 02.41 M B7 III Taurus 


Alcyone is the brightest star in the Pleiades star cluster, with a luminosity of 
about 350 times that of the Sun. 

Maia 20Tauri 03'’45.8 m +24°22' Oct-Nov-Dec 

3.87 m - 1.344 M B8III Taurus 


This is yet another lovely blue star in the Pleiades cluster. Maia has a luminosity 
about eight times that of the Sun. 


30 


Astrophysics is Easy 


EtaSagitai eSgr 18 , '24.2 ra -34°23' May-Jun-Jul 

1.79 m — 1.44 M B9.5III Sagittarius 

This is a brilliant orange star at a distance of 125 l.y. with a luminosity 250 times 
that of the Sun’s. 

NuDraconis 1 v'Dra 17 ,, 32.2 m +55°11' May-Jun-Jul 

4.89 m 2.48 M Am Draco 

This is a classic double star system visible through binoculars or a small telescope. 
The stars are nearly identical in magnitude and stellar class and have a lovely 
white color. 


Alhena yGem 06 h 37.7 m +16°23' Nov-Dec-Jan 

1.93 m — 0.60 M AO IV Gemini 

The star is relatively close at 58 l.y. with a luminosity 160 times that of the Sun. 

Castor a Gem 07 ,I 34.6 m +31°53' Dec-Jan-Feb 

1.43 m 0.94 M A1 V Gemini 

Castor is a part of the famous multiple star system and a fainter brother to Pollux. 
The visible magnitude stated here is the result of combining the magnitudes of 
the two brighter components of the system, 1.9 and 2.9. 

Deneb aCyg 20' ! 41.3” ! +45°17' Jul-Aug-Sep 

1.25 v m — 8.73 36 M A2 I Cygnus 

This is the faintest star of the Summer Triangle (the others being Altair and 
Vega). Deneb is a supergiant star with definite pale-blue color. It is the prototype 
of a class of pulsating variable stars. 

Denebola [iLeo ll ,! 49.1 m +14°34' Feb-Mar-Apr 

2.14 v m 1.92 M A3 V Leo 

This star, along with several of its companion stars, is visible through a variety 
of instruments. The star has just recently been designated a variable. 

Delta Leonis 8 Leo ll , '14.1 m +20°31' Feb-Mar-Apr 

2.56 m 1.32 M A4V Leo 

Also called Zozma, Delta Leonis lies at a distance of 80 l.y. with a luminosity 50 
times that of the Sun. 

RasAlhague a Oph 17 , '34.9 m +12°34' May-Jun-Jul 

2.08 m 1.30 M A5 III Ophiucus 

This is an interesting star for several reasons. It shows the same motions through 
space as several other stars in the Ursa Major Group. It also shows interstellar 
absorption lines in its spectrum. Finally, its measurements show an oscillation 
(or wobble) in proper motion, which would indicate an unseen companion star. 


Tools of the Trade 


31 


2 Mon HD 40536 05 ft 59.1 m -09°33' Nov-Dec-Jan 

5.01m 0.02 M A6 Monoceros 

The star lies at a distance of over 1900 l.y. with a luminosity 5000 times that of 
the Sun. 

Alderamin a Cep 21 ,1 18.6 m +62°35' Jul-Aug-Sep 

2.45 m 1.58 M A7IV Cepheus 

This is a rapidly rotating star, which results in its spectral lines becoming broad 
and less clear. It also has the dubious distinction of becoming the Pole Star in 
7500 A.D. 

Gamma Herculis yHer 16 h 21.8 m +19°09' Apr-May-Jun 

3.74 m — 0.15 M A9 III Hercules 

This is an optical double system, lying at a distance of 144 l.y., with a luminosity 
46 times that of the Sun. 

Canopus a Car 06 ,1 23.9 m — 52°41' Nov-Dec-Jan 

—0.62 m — 5.53M F0 1 Carina 

This is the second brightest star in the sky. Its color is often reported as orange 
or yellow, as it is usually seen low down in the sky, and hence it is more likely 
to be affected by the atmosphere. Its true color is white. 

b Velorum HD 74180 08 h 40.6 m -46°39' Dec-Jan-Feb 

3.84 m -6.12 M F3 I Vela 

This star is unremarkable except that its luminosity is estimated to be 180,000 
times that of the Sun! 

Zubenelgenubi a 'Lib 14 ,! 50.7 m —15° 60' Apr-May-Jun 

5.15m 3.28M F4 IV Libra 

An easily resolvable double star, ot 1 is a spectroscopic binary. The colors are a 
nice faint yellow and pale blue. 

Algenib a Per 03 A 24.3 m +49°52' Oct-Nov-Dec 

1.79 m — 4.5 M F5 I Perseus 

The star lies within Melotte 20, a loosely bound stellar association, also known as 
Perseus OB-3 or Alpha Persei Association. About 75 stars with magnitude 10 or 
below are contained within this group. These are stellar infants, 50 million years 
old, lying 550 l.y. away. The metallic lines increase through the F class, especially 
the H and K lines of ionized calcium. 

Polaris aUMi 02 , ’31.8 m +89°16' Sep-Oct-Nov 

1.97 v m — 3.64 M F7 I Ursa Minor 

Polaris is an interesting and famous star, yet it is only the 49th brightest star in 
the sky. It is a Cepheid Variable type II star (the W Virginis class), and a binary 
star (the companion reported as being pale blue). The star is expected to move 
closest to the celestial pole in 2102 A.D. 


32 


Astrophysics is Easy 


pVir HD 102870 ll' ! 50.7 m +01°46' Feb-Mar-Apr 

3.59m 3.40M F8 V Virgo 

A close star at 34 l.y., (3 Vir is just three times as luminous as the Sun. 

Sadal Suud P Aqr 21 h 31.6 m — 05°34' Jul-Aug-Sep 

2.90m — 3.47M GO I Aquarius 

A giant star and a close twin to a Aqr, Sadal Suud lies at a distance of 990 l.y. 
and is 5000 times more luminous than the Sun. 

Sadal Melik a Aqr 22 ,! 05.8 m -00° 19' Jul-Aug-Sep 

2.95m — 3.88M G2 1 Aquarius 

Although it has the same spectral class and surface temperature as that of the 
Sun, a Aqr is a giant star, whereas the Sun is a main-sequence star. 

Ras Algethi a 2 Her ll h \4.7 m +14°23' May-Jun-Jul 
5.37m 0.03M G5 III Hercules 

This is a beautiful double star with colors of ruddy orange and bluish green. 
The spectral class refers to the primary star of a 2 Her, which is a spectroscopic 
double, and thus visually inseparable with any telescope. 

Algeiba y 2 Leo 10 ,! 19.9 m +19°50' Jan-Feb-Mar 

3.64m 0.72M G7 III Leo 

Algeiba is a famous double star. Most observers report an orange or yellow color, 
but some report the G7 star as greenish. 

P LMi HD 90537 10 ,! 27.8 m +36° 42' Jan-Feb-Mar 

4.20m 0.9M G8 III Leo Minor 

A constellation in which no star is given the classification a, (3 LMi has the 
misfortune not being the brightest star in the constellation, the honor of which 
goes to 46 LMi. 

P Cet HD 4128 00 , '43.6 m -17°59' Sep-Oct-Nov 

2.04m -0.30 M G9.5 III Cetus 

This star lies at a distance of 60 l.y. with luminosity 42 times that of the Sun. 

Gienah e Cyg 20 ,! 46.2 m +33°58' Jul-Aug-Sep 

2.48m 0.76M K0 III Cygnus 

Marking the eastern arm of the Northern Cross, Gienah is a spectroscopic binary. 
In the K-class stars, the metallic lines are becoming more prominent than the 
hydrogen lines. 

v 2 CMa HD 47205 06 h 36.7 m -19°15' Nov-Dec-Jan 

3.95m 2.46M K1 III Canis Major 


Tools of the Trade 


33 


This star lies at a distance of 60 l.y. with luminosity seven times that of the Sun. 

Enif e Peg 2l h 44.2 m +09°52' Jul-Aug-Sep 

2.38 v m — 4.19 M K2 I Pegasus 

This star lies at a distance of 740 l.y. with luminosity 7,450 times that of the Sun. 
The two faint stars in the same field of view have been mistakenly classified as 
companions, but on analysis they have been proved to be stars in the line of 
sight. 


Almach y'And 02 h 03.9 m +42°20' Sep-Oct-Nov 

2.33m — 2.86 M K3 III Andromeda 

This is a famous binary star whose colors are golden and blue, although some 
observers see orange and greenish-blue. Nevertheless, the fainter companion is 
hot enough to show a truly blue color. It is also a binary in its own right, but 
not observable through amateur instruments. 

£ 2 Sco HD 152334 16''54.6 m -42°22' May-Jun-Jul 

3.62 m 0.3 M K4 III Scorpius 

The brighter of the two stars in this naked-eye optical double star system, the 
orange supergiant star contrasts nicely with its slightly fainter blue supergiant 
companion. 

v'Boo HD 138481 15 h 30.9 m +40°50' Apr-May-Jun 

5.04m — 2.10 M K5 III Bootes 

The star lies at a distance of 385 l.y. and has a luminosity 104 times that of the 
Sun (see also Aldebaran ). 

Mirach PAnd 01 h 09.7 m +35°37' Sep-Oct-Nov 

2.07m — 1.86 M M0 III Andromeda 

In this stellar class, the bands of titanium oxide are strengthening. This red giant 
star is suspected to be slightly variable, like so many other stars of the same type. 
In the field of view is the Galaxy NGC 404. 

Antares a Sco 16 , ’29.4 m — 26°26' Apr-May-Jun 

1.06 v m — 5.28 M Ml I Scorpio 

This giant star, measuring 600 times the diameter of the Sun, has a glorious fiery 
red color, contrasting nicely with its fainter green companion. 

Scheat p Peg 23 ,1 03.8'" +28° 45' Aug-Sep-Oct 

2.44 v m — 1.49 M M2 II Pegasus 

Marking the north-western corner of the Square of Pegasus, Scheat is a red, 
irregular variable star. It was noted for having been one of the first stars to have 


34 


Astrophysics is Easy 


its diameter (0.021") measured by the technique of interferometry. Being a 
variable star, its size oscillates to a maximum diameter of 160 times that of the 
Sun. 


Eta Persei T| Per 02 h 50.7' n +55°54' Oct-Nov-Dec 

3.77m — 4.28M M3 1 Perseus 


This yellowish star is an easily resolved double-star system. The color contrasts 
nicely with its blue companion. 


Gacrux y A Cruris 12 h 3l.2' n -57°07' Feb-Mar-Apr 

1.59m — 0.56M M4 III Crux 

The top star of the Southern Cross, Gacrux is a giant star. y A and y B do not form 
a true binary since they are apparently moving in different directions. 


Ras Algethi a 1 Her 17 ,1 14.6 m +14°23' May-Jun-Jul 

3.03 v m — 2.32 M M5 II Hercules 


A fine double-star system. The M5 semi-regular star is an orange supergiant, 
which contrasts with its companion, a blue-green giant. It must be pointed out 
that this double star can be resolved only with a telescope (not with binoc- 
ulars), as the two stars are less than 5" apart. The change in brightness may be 
attributed to actual physical changes to the star as it increases and decreases in 
diameter. 


Mira (at maximum) o Cet 02 ,, 19.3 m -02°59' Sep-Oct-Nov 

2.00 v m — 3.54 M M5 Cetus 

For details on Mira, see the section “Long Period Variables.” 

Mira (at minimum) oCet 02' 1 19.3 m -02°59' Sep-Oct-Nov 

10 v m — 0.5 M M9 Cetus 

For details on Mira, see the section “Long Period Variables.” 

0 Apodis HD 122250 14 ,! 05.3 m -76° 48' Mar-Apr-May 

5.69 v m -0.67 M M6.5 III Apus 


This is a semi-regular variable with a period of 119 days in the range of 5th to 
nearly 8th magnitude. The titanium bands are now at their strongest. 


Tools of the Trade 


35 


1 .9 The Hertzsprung-Russell 
Diagram 


We have already covered many topics in our description of a star’s basic charac- 
teristics, such as its mass, radius, spectral type, and temperature. Let us now put 
all these parameters together to get a picture of how a star evolves. It is often 
quite useful in many subjects to represent the data about a group of objects 
in the form of a graph. Many of us are familiar with graphs or have seen, for 
example, one that shows height as a function of age, or temperature as a function 
of time. A similar approach has been pursued to study the characteristics of stars. 
A graph that is used universally is called The Hertzsprung-Russell Diagram. It is, 
without a doubt, one of the most important and useful diagrams in the study of 
astronomy. 

In 1911, the Danish astronomer Ejnar Hertzsprung plotted the absolute 
magnitude of stars (a measure of their luminosities) against their colors 
(a measure of their temperature). Later, in 1913, the American astronomer 
Henry Norris Russell independently plotted spectral types (another way to 
measure temperature) against absolute magnitude. They both realized that 
certain unsuspected patterns began to emerge, and furthermore, an under- 
standing of these patterns was crucial to the study of stars. In recognition 
of the pioneering work of these astronomers, the graph was known as the 
Hertzsprung-Russell Diagram, or H-R diagram. Figure 1.4 is a typical H-R 
diagram. Each dot on the diagram represents a star whose properties, such as 
spectral type and luminosity, have been determined. Note the key features of the 
diagram: 

• The horizontal axis represents stellar temperature or, equivalently, the spectral 
type. 

• The temperature increases from right to left. This is because Hertzsprung and 
Russell originally based their diagram on the spectral sequence OBAFGKM, 
where hot O-type stars are on the left and cool M-type stars on the right. 

• The vertical axis represents stellar luminosity measured in the unit of Sun’s 
luminosity, L 0 . 

• The luminosities cover a wide range, so the diagram makes use of the 
logarithmic scale, whereby each tick mark on the vertical axis represents a 
luminosity 10 times larger than the prior one. 

• Each dot on the H-R diagram represents the spectral type and luminosity of 
a single star. For example, the dot representing the Sun corresponds to its 
spectral type G2 with luminosity L 0 = 1. 

Note that because luminosity increases upward in the diagram and surface 
temperature increases leftward, stars near the upper left corner are hot and 
luminous. Similarly, stars near the upper right corner are cool and luminous; 
stars near the lower right corner are cool and dim; and finally stars near the 
lower left corner are hot and dim. 


36 


Astrophysics is Easy 


Surface temperature (K) 

25,000 10,000 S000 6000 5000 4000 3000 



Spectral type 


Figure 1.4. The Hertzsprung-Russell Diagram. Luminosity is plotted against spectral type for a 
selection of stars. Some of the brighter stars are shown. Each dot represents a star whose 
spectral type and luminosity have been determined. Note how the data are grouped in just a 
few regions, indicating a correlation. The main sequence is the continuous blue line. Surface 
temperature and absolute magnitude are also shown. 



Tools of the Trade 


37 


1.10 The H-R Diagram 
and Stellar Radius 


The H-R diagram can directly provide important information about the radius 
of stars, because the luminosity of a star depends on both its surface temper- 
ature and surface area, or radius. You may recall that the surface temper- 
ature determines the amount of power emitted by the star per unit area. 
Thus, a higher temperature means a greater power output per unit area. So, 
if two stars have the same temperature, the larger star may be more luminous 
than the other. Stellar radii must perforce increase as we go from the high- 
temperature, low-luminosity corner on the lower left of the H-R diagram to 
the low-temperature, high-luminosity upper right corner. This is shown in 
Figure 1.5. 

The first thing to notice on the H-R diagram is that the data points (or stars) 
are not scattered at random but appear to fall into distinct regions. This would 
imply that surface temperature (or spectral type) and luminosity are related! The 
several groupings can be described as thus: 


• The band that stretches diagonally across the H-R diagram is called the Main 
Sequence, and it represents about 90% of the stars in the night sky. It extends 
from hot and luminous blue stars in the upper left corner to cool and dim 
red stars in the bottom right. Any star located in this part of the H-R diagram 
is called a main-sequence star. Note that the Sun is a main-sequence star 
(spectral type G2, absolute magnitude +4.8, luminosity 1 L 0 ). We shall see 
later in the book that stars on the main sequence are undergoing hydrogen- 
burning (thermonuclear fusion, which converts hydrogen to helium) in their 
cores. 

• Stars in the upper right are called giants. These stars are both cool and 
luminous. Recall from an earlier section that we discussed the Stefan- 
Boltzmann Law, which states that a cool star will radiate much less energy 
per surface area than a hot star. Hence, for these stars to appear as luminous 
as they look, they must be immense, and so they are called giants. They may 
be anywhere from 10 to 100 times as big as the Sun. Figure 1.5 shows this, 
where stellar radii have been added to the H-R diagram. Most giant stars 
are about 100 to 1000 times more luminous than the Sun and have tempera- 
tures of 3000 to 6000 K. Many of the cooler members of this class are reddish 
and have temperatures of 3000 to 4000 K — these are often referred to as red 
giants. Some examples of red giants are Arcturus in Bootes and Aldebaran in 
Taurus. 

• At the extreme upper right corner are a few stars that are even bigger than the 
giants. These are the supergiants, which have radii up to 1000 R Q . Giants and 
supergiants make up about 1% of stars in the night sky. Antares in Scorpius and 
Betelgeuse in Orion are two fine examples of supergiant stars. Nuclear fusion 
taking place in supergiant stars is significantly different in both character and 
position than the reactions taking place in the stars on the main sequence. 



38 


Astrophysics is Easy 


io 6 


10 “ 


10 - 


Luminosity (Ls) 


io-- 


10 4 



1 - 


40,000 20,000 


10,000 5000 

Surface temperature (K) 


3000 


Figure 1.5. Size of stars on an H-R diagram. Stellar luminosity against surface temperature. 
The dashed diagonal lines indicate stars of different radii. At a given radius, the surface 
temperature increases (moving from right to left), and luminosity increases. Notice the main 
sequence and the Sun's position on it. A very average star. 


• Stars in the lower left of the H-R diagram are much smaller in radius and 
appear white. These are the white dwarf stars. As we see from the H-R diagram, 
they are hot stars with low luminosities; therefore, they must be small and 
hence the name dwarf stars. They are faint stars, and so they can be seen 
only with telescopes. They are approximately the same size as the Earth. There 
are no nuclear reactions within white dwarfs; rather, they are the still-glowing 
remnants of giant stars. White dwarfs account for about 9% of stars in the 
night sky. 


Tools of the Trade 


39 


1.11 The H-R Diagram 
and Stellar Luminosity 


The temperature of a star determines which spectral lines are most prominent 
in its spectrum. Therefore, classifying a star by its spectral type is essentially the 
same as by its temperature. A quick look at an H-R diagram will reveal that stars 
can have similar temperatures but in fact very different luminosities. 

Consider this example: a white dwarf star may have a temperature of 7000 K; 
so do a main-sequence star, a giant, and a supergiant. It all depends on its 
luminosity. Therefore, by examining a star’s spectral lines, one can determine 
which category the star belongs to. A rule of thumb (for stars of spectral types B 
through F) is: the more luminous the star, the narrower the lines of hydrogen. 
The theory behind the phenomenon is quite complex, but suffice to say that these 
measurable differences in spectra are due to differences in stars’ atmospheres 
where absorption lines are produced. The density and pressure of hot gases in 
the atmosphere affect the absorption lines and hydrogen in particular. If the 
pressure and density are high, hydrogen atoms collide more frequently, and they 
interact with other atoms in the gas. The collisions cause the energy levels in 
hydrogen atoms to shift, resulting in broadened hydrogen spectral lines. 

In a giant luminous star, the atmosphere will have a very low pressure 
and density because the star’s mass is spread over such an enormous volume. 
Therefore, the atoms (and ions) are relatively far apart. This means that collisions 
between atoms are far less frequent, which produces narrow hydrogen lines. In 
a main-sequence star, the atmosphere is denser than a giant or supergiant, with 
collisions occurring more frequently, thereby producing broader hydrogen lines. 

In an earlier section on “Stellar Classification,” we saw that we can ascribe 
to a star a luminosity class. We can use it here to describe the region of the 
H-R diagram where a star of a particular luminosity will fall. This is shown in 
Figure 1.6. 

Knowing both the spectral type and luminosity of a star would help an 
astronomer to instantly know where on the main sequence it lies. For instance, a 
G2 V star is a main-sequence star with a luminosity of 1 L 0 and a surface temper- 
ature of about 5700 K. In a similar vein, Aldebaran is a K5 III star, which means 
that it is a red giant star with a luminosity of 375 L Q and a surface temperature 
of about 4000 K. 


1.12 The H-R Diagram 
and Stellar Mass 


The most common trait of main-sequence stars is that, just like the Sun, they 
undergo nuclear fusion at their cores to convert hydrogen to helium. Since most 
stars spend much part of their lives doing this, it naturally follows that a majority 
of stars spend their time somewhere on the main sequence. Even a cursory look 




40 


Astrophysics is Easy 



Spectral type 

Figure 1.6. Luminosity classes. Dividing the H-R diagram according to luminosity classes 
allows distinctions to be made between giant and supergiant stars. 


at the H-R diagram can tell you that an enormous range of luminosities and 
temperatures are covered. 

The question that may arise is, why such a large range? 

Astronomers have determined the masses of stars using binary star systems, 
and they discovered that a star’s mass increases as it moves upward along the 
main sequence (Figure 1.7). The 0-type stars, which are hot and luminous stars, 
at the upper part of the diagram can have masses as high as 100 times that of 
the Sun — 100M o . At the other end of the main sequence, the cool and faint 



Tools of the Trade 


41 


io c 


10- 1 


10 J 


Luminosity (L s ) 


Iff- 


10 J 


60 
* 37 


1 24 


17 


. 7.7 


3.9 


2.8 


2.1 


1.6 


' 1.4 


Sun 
• 0.78 


* 0.68 
• 0.50 
• 0.38 


• 0.22 


40,000 20,000 10,000 5000 

< Surface temperaure (K) 


3000 


Figure 1.7. Mass and the main sequence. Each filled-in circle is a main-sequence star. The 
number is the star's mass in solar masses (M G ). As you move up the main sequence (from lower 
right to upper left), the mass, luminosity, and temperature increase. 


stars have masses as low as 0.1 time that of the Sun — O.1M 0 . 37 This orderly 
distribution of stellar masses along the main sequence tells us that mass is the 
most important attribute of a hydrogen-burning star. Mass has a direct effect on 
a star’s luminosity because the weight of a star’s outer layers will determine how 
fast the hydrogen-to-helium nuclear reaction will proceed in the core. A 10 M Q 
star on the main sequence will be more than 1000 times more luminous than the 
Sun (i.e., 1000 L Q ). 



42 


Astrophysics is Easy 


However, the mass-surface temperature relationship is just a little more subtle 
than the preceding paragraph indicated. Generally, very high luminous stars must 
either be very large or have a very high temperature, or even a combination of 
both. Stars on the top left corner of the main sequence are some thousands of 
times more luminous than the Sun, but they are only about 10 times larger than 
the Sun. Therefore, their surface temperatures must be significantly hotter than 
that of the Sun to account for such high luminosities. Bearing this relationship 
in mind, we can now say that main-sequence stars that are more massive than 
the Sun must have correspondingly higher temperatures, while those with lower 
masses must have lower surface temperatures. Thus, you can now understand 
why the main sequence on the H-R diagram goes diagonally from upper left to 
the lower right. 

The H-R diagram is one of the most fundamental tools in astronomy. We will 
use it throughout the remainder of this book, as it provides a means to determine 
the many paths that stars take during their lives — from star birth to star death. 


Notes 


1. One parsec is equal to 3.26 light years, 3.09 x 10 13 km, or 206265 AU. 1 AU 
is 149,597,870 km. 

2. Nearly 200 previously unobserved stars were discovered, the nearest about 
18 l.y. away. In addition, several hundred stars originally believed to be 
within 75 l.y. were in fact found to be much farther away. 

3. The most famous Cepheid variable star is Polaris, the North Star. It varies its 
visual brightness by about 10% in just under 4 days. Recent data show that 
the variability is decreasing, and the star may, at some point in the future, 
cease to pulsate. We shall discuss Polaris and other important variable stars 
in detail in a later section. 

4. We shall discuss the meaning of the term luminosity later. For the time 
being, think of it as the star’s brightness. 

5. The Period-Luminosity relationship was discovered by Henrietta Leavitt 
in 1908 while working at the Harvard College Observatory. She studied 
photographs of the Magellanic Clouds and found more than 1700 variable 
stars. 

6. The relationship between the apparent brightness of a star and its intrinsic 
brightness will be discussed in the next section. 

7. This signifies that the star is in fact part of a double star system, and the 
distance quoted is for components A and B. 

8. The star, and thus the magnitude, is variable. 

9. Most of the nearest stars are very faint, so only the brighter ones will be 
mentioned here. Exceptions to this will be made, however, if the object 
has an important role in astronomy. A companion book to this one — Field 
Guide to the Deep Sky Objects — provides much more information and detail 
regarding the nearest stars. Furthermore, the Field Guide addresses many 
techniques to enhance your observational skills, such as dark adaption, 
averted vision, etc. 



Tools of the Trade 


43 


10. The proper motion of a star is its apparent motion across the sky. 

11. The HD signifies it is the 217987th object in the Henry Draper Catalogue. 

12. One watt is equal to 1 joule per second. The Sun’s luminosity is 3.86 x 
10 26 W. It is often designated by the symbol L 0 . 

13. The scientific term for apparent brightness is flux. 

14. Observers have reported that, under excellent conditions and with very 
dark skies, objects down to magnitude 8 can be seen with the naked eye. 

15. Many stars are variable, so the value for their apparent magnitude will 
change. The suffix v indicates a variable star, and the value given is the 
mean value. 

16. It shouldn’t come as any surprise to you to learn that there are several other 
magnitude definitions that rely on a star’s brightness when observed at a 
different wavelength — the U, B, and V system. There is also a scale based on 
photographic plates, the photographic magnitude, m pg , and the photovisual 
magnitude, m pv . Finally, there is the bolometric magnitude, m B0L , which is 
a measure of all the radiation emitted from an object. 

17. The eye does not recognize color at low light levels. This is why at night, 
with the naked eye, we see only shades of grey, white, and black. 

18. The most important factor determining the color of a star you see is 
you — the observer! It is purely a matter of physiological and psychological 
influences. What one observer describes as a blue star, another may describe 
as a white star; or one may see an orange star, while another observes 
the same star as yellow. You might even observe a star to have different 
color when using different telescopes or magnifications, and atmospheric 
conditions certainly have a role to play. 

19. From here on, when I mention temperature, I am referring to the surface 
temperature, unless indicated otherwise. 

20. This star is the brightest in the night sky. It is, of course, Sirius. 

21. This is the most famous irregular variable star, Mira. 

22. Remember that a star’s color is observer-dependent! What one person sees 
as yellow, another sees as white. Do not be surprised if you see a different 
color to that mentioned. 

23. The real temperature of the star is still undetermined. 

24. A few stars, such as Betelgeuse, have had their radii determined by a 
technique known as interferometry. For the vast majority of stars, the 
technique is not applicable, either due to distance or faintness. 

25. To be accurate, the law refers to a black-body, which is something that 
emits thermal radiation. Thus, thermal radiation is blackbody radiation. It 
can be applied to a star because, to all intents and purposes, a star’s surface 
behaves like a black-body. 

26. No doubt some of you are already asking, “Where is the surface of a star? 
A star is made of gas.” Fear not... all will be revealed in later chapters. 

27. There is some uncertainty about this value. 

28. Astronomers call every element other than hydrogen and helium a metal. 
It’s odd, I agree, but don’t worry about it — just accept it. 

29. We can only easily deduce whether an object is moving away from us 
or toward us. To measure if it is moving laterally to us requires some 
complicated measurements. 


44 Astrophysics is Easy 

30. Some spectroscopes place the prism or grating in front of the telescope, and 
thus the light from every star in the field of view is analyzed simultaneously. 
This is called an objective spectroscope. The drawback is the considerable 
loss of detail (i.e., information about the stars), but initial measurements 
can be made. 

31. The transitions shown are only a few of the many that occur. 

32. The reason that stars follow the order OBAFGKM was discovered by a 
brilliant astronomer, Cecilia Payne-Gaposchkin. She found that all stars are 
made primarily of hydrogen and helium and that a star’s surface temper- 
ature determines the strength of its spectral lines. For instance, O stars have 
weak hydrogen lines because, due to their high temperature, nearly all the 
hydrogen is ionized. Thus, without an electron to “jump” between energy 
levels, ionized hydrogen can neither emit nor absorb light. On the other 
hand, M stars are cool enough for molecules to form, resulting in strong 
molecular absorption lines. 

33. As we shall see later, these are stars with very low temperatures — 1900 
to 1500 K. Many astronomers now believe these are the infamous brown 
dwarfs. 

34. These can be further classified into la and lb, with la the brighter. 

35. Usually only the classes O, A, B, F, G, K, and M are listed. The other classes 
are used and defined as and when they are needed. 

36. This value is in question. The data are awaiting reassessment. 

37. Over the past several years, astronomers have discovered that the low-mass, 
faint M-type dwarf stars are far more numerous than other star types. We 
have just not been able to see them up until now. 


CHAPTER TWO 


The Interstellar 
Medium 


2.1 Introduction 


When we look up into the night sky, we see stars, and not much else. So we 
get the impression that between the stars, space is empty. There does not seem 
to be any sort of material that lies between one star and another. At the same 
time, we know intuitively that this cannot be true, for if space were empty, 
from what did stars form? This then leads us to the conclusion that perhaps 
space is not quite so empty, but filled with some sort of material that, to our 
eyes, is all but invisible yet is responsible for providing the source material for 
stars. 

In fact, space is anything but empty; it is filled with gas and dust. This is 
known as the Interstellar Medium (ISM). The ISM is made up of gas (mainly 
hydrogen) and dust (which accounts for about 1% of the mass of gas). The dust, 
not to be confused with dust on Earth, consists of other elements that are not 
hydrogen, such as carbon, silicon, and so on, and their compounds, CO, HCN, 
and so on. 

The material that makes up the ISM is not spread evenly throughout 
space; there are regions that are dense and regions that are not so dense. Similarly, 
there are areas of the ISM that are hot and other areas that are cooler. Thus, 
the two most important parameters concerning the ISM are the temperature 
and a quantity we call the number density ( n ). The latter is just the number 
of particles per unit volume (per cubic meter), and it can be individual atoms, 
neutral, ionized, combined in molecules, or a mixture of all four. Because there is 


45 


46 


Astrophysics is Easy 


far more hydrogen in the ISM than anything else, we can say, to a good approx- 
imation, that the particle density ( n ) is the number of hydrogen atoms per cubic 
meter, and this we call n H 

The important point to realize is the enormous range of temperatures and 
number densities that occur in the ISM. There can be as few as 100 particles per 
cubic meter {n — 100 m~ 3 ) to about 10 17 per cubic meter ( n = 10 17 m“ 3 ). Similarly, 
the temperature can be as low as 10 K and as high as a few million K. To get a 
feel for these ranges, look at Figure 2.1 It shows the ranges of temperature and 
number density, and the names we give to the correspondingly different regions 
in the ISM. 

Let us look at this diagram in more detail; what we call the intercloud medium, 
whether hot or warm, actually accounts for most of the ISM. The interesting thing 
is that all other regions of the ISM are located within the intercloud medium. 
The regions are: 

Hot intercloud medium — this is widespread and, although hot, of extremely 
low density, and consisting mainly of ionized hydrogen. Fortunately for amateur 


1,000,000,000,000,000 _ 

Circumstellar 

shells 

Inner regions 



1,000,000,000,00 - 

Dense 

clouds 



Number of 




particles per 1 ,000,000,000 - 
cubic meter, (ft) 


Planetary 

nebulae 


10,000,000 - 

Diffuse clouds 

HII 

regions 

Supernova 



Warm 

intercloud 

medium 

remnants 

1000 - 



Hot 

intercloud 

medium 


1 100 10,000 1,000,000 

(K) (K) (K) (K) 

Temperature (K) 

Cold Warm Hot 


Figure 2.1. Regions in the interstellar medium. 



The Interstellar Medium 


47 


astronomers, it does not obscure our view of space, as we can see through it. For 
similar reasons, the warm intercloud medium is also transparent. 

All other regions on the diagram (and, thus, in the ISM) present a much 
more visual aspect and so are important to us as observers. They can be divided 
into two groups: the regions in the ISM that are concerned with star formation, 
namely the diffuse and dense clouds and the HII regions, 2 and those that deal with 
star death — planetary nebulae, supernova remnants, and circumstellar shells. 

We shall discuss all these regions in considerable detail in this and later 
chapters because they are objects of interest to the amateur astronomer. . .all, that 
is, except for the diffuse clouds, as these are transparent to visible light. There 
are, however, methods to allow observations of these clouds; radio astronomy 
to measure the hydrogen 21 cm line, microwave telescopes to measure the CO 
molecule, and infrared telescopes to measure the far infrared emission of the 
dust. 

As an astronomer, you will have already observed the interstellar medium, 
but perhaps without realizing it. As previously mentioned, the ISM is composed 
of gas (mainly hydrogen), 3 and dust, so it is, from an observer’s viewpoint, 
invisible; however, there are places in the Galaxy where certain conditions tend to 
aggregate the material, and these denser-than-average regions are indeed visible 
to the amateur astronomer. We know them as nebulae. 


2.2 Nebulae 


Nebulae are actually disparate in nature, even though many of them have a 
rather similar appearance. They are associated with the areas of star formation, 4 
cover several aspects of a star’s life, and end with the process of star death. 
This section will just cover three main types of nebulae: emission, reflecting, and 
dark, all associated with the birth of a star. In addition, we are fortunate from 
an observing point of view, because these objects abound in the night sky, and 
some are spectacular objects indeed. 


2.3 Emission Nebulae 


These clouds of gas are associated with very hot O- and B-type stars, which 
produce immense amounts of ultraviolet radiation. They typically have masses of 
about 100 to 10,000 solar masses. This huge mass, however, is usually spread over 
a correspondingly large area (possibly a few l.y. across), so the actual density of 
the gas is extremely low (maybe only a few thousand hydrogen atoms per cubic 
centimeter). Usually, these very luminous stars are actually born within and from 
the material of the clouds, and so many emission nebulae are “stellar nurseries.” 
Radiation from the stars causes the gas (usually hydrogen) to undergo a process 
called fluorescence, and it is this process that is responsible for the glow observed 
from the gas clouds. 






48 


Astrophysics is Easy 


The energy provided by ultraviolet radiation from the young and hot stars 
ionizes the hydrogen. In other words, energy — in this case, in the form of ultra- 
violet radiation — is absorbed by the atom and transferred to an electron that 
is sitting comfortably in what is called an energy level or orbital shell. 5 Having 
gained extra energy, the electron can leave the energy level it is in, and in some 
instances actually break free from the atom. When an atom loses an electron, 
the process is called ionization. 

If electrons are broken free from their parent atoms, the hydrogen cloud 
will contain some hydrogen atoms without electrons — ionized hydrogen (also 
known as protons), and a corresponding number of free electrons. Eventually, 6 
the electrons recombine with the atoms, but an electron cannot just settle down 
back to the state it was originally in before it absorbed the extra energy — it has 
to lose the extra energy that the ultraviolet imparted. For this to happen, the 
electron moves down the atomic energy levels until it reaches its original level, 
losing energy as it goes. In hydrogen (the most common gas in the nebula), an 
electron moving down from the third energy level to the second emits a photon 
of light at 656.3 nm (see Section 1.7). 

This is the origin of the famous “hydrogen alpha line,” usually written as 
H-alpha, and pronounced “aitch alpha.” It is a lovely red-pink color and is 
responsible for all the pink and red glowing gas clouds seen in photographs of 
emission nebulae. 7 

When electrons move down from other energy levels within the atom, other 
specific wavelengths of light are emitted. For instance, when an electron moves 
from the second level to the first, it emits a photon in the ultraviolet part of the 
spectrum. This particular wavelength is called the Lyman alpha line of hydrogen, 
which is in the ultraviolet part of the spectrum. 

It is this process of atoms’ absorbing radiation to ionize a gas, with electrons 
subsequently cascading down the energy levels of an atom, that is responsible for 
nearly all of the light we see from emission nebulae. If a gas cloud is particularly 
dense, the oxygen gas in it may be ionized, and the resulting recombination 
of the electron and atom produces the doubly ionized lines, at wavelengths of 
495.9 and 500.7 nm. 8 

Emission nebulae are sometimes called HII regions, pronounced “aitch two.” 
This astrophysical term refers to hydrogen that has lost one electron by 
ionization. The term HI, or “aitch one,” refers to hydrogen that is unaffected by 
any radiation (i.e., neutral hydrogen). The doubly ionized oxygen line mentioned 
above is termed OIII (“oh three”); the “doubly” means that two of the outermost 
electrons have been lost from the atom by ionization. 9 

The shape of an emission nebula is dependent on several factors: the amount 
of radiation available, the density of the gas cloud, and the amount of gas 
available for ionization. When there is a significant amount of radiation, coupled 
with a small and low-density cloud, then all of the cloud will likely be ionized, 
and thus the resulting HII region will be of an irregular shape — just the shape 
of the cloud itself. If the cloud of gas is large and dense, however, then the 
radiation can penetrate only to a certain distance before it is used up — that 
is, there is only a fixed amount of radiation available for ionization. In this 
case, the HII region will be a sphere, 10 often surrounded by the remaining 
gas cloud, which is not fluorescing. Many of the emission regions that are 


The Interstellar Medium 


49 


irregular in shape include M42 (the Orion Nebula), M8 (the Lagoon Nebula), 
and M17 in Sagittarius. Two of those exhibiting a circular shape, and thus 
are in fact spherical, are M20 (the Trifid Nebula) and NGC 2237 (the Rosette 
Nebula). 

After a suitable period of time, usually several million years, the group of 
young O- and B-type stars located at the center of the nebulae will be producing 
so much radiation that they can in effect sweep away the residual gas and dust 
clouds that surround them. This produces a “bubble” of clear space surrounding 
the cluster of stars. Several emission regions show this. For example, NGC 6276 
and M 78 show the star cluster residing in a circular clear area within the larger 
emission nebula. 

Let us now look at a few examples of the brighter emission nebulae. Note, 
however, that from an observational viewpoint, many of the emission nebulae 
are faint and have a low surface brightness, making them not exactly difficult 
objects to observe but rather featureless and indistinct (though in some instances 
the brighter nebulae do show several easily seen features). Therefore, clear nights 
and clean optics are a high priority. 

The photographic brightness (as seen on the photographic plates from the 
Palomar Observatory Sky Survey (POSS) is assigned a value from 1 through 6; 
those nebulae rated at 1 are just barely detectable on the plate, while those 
quoted at a value of 6 are easily seen on the photographic plate, and this number 
is just the measure of the difficulty (or ease) of observation, and is given the 
symbol #. The size of an object is also given (in arcseconds), and is indicated by 
the symbol ©. Where a value of © is given as x^y, the object is approximately 
x arcseconds long by y arcseconds wide. 


2.3.1 Brightest Emission Nebulae 

Gum 4 NGC 2359 07 h 18.6 m -13°12' Dec-Jan-Feb 

*112-5 ©9^46' Canis Major 

Also known as the Duck Nebula, this is a bright emission consisting of two 
patches of nebulosity, with the northern patch being the larger and less dense. Try 
using an OIII filter to improve the appearance of the emission nebula, showing 
the delicate filamentary nature. 

Messier 20 NGC 6514 18 h 02.3 ra -23°02' May-Jun-Jul 

#1-5 ©20^20' Sagittarius 

Also known as the Trifid Nebula, this emission nebula can be glimpsed as 
a small hazy patch of nebulosity and is easy to see, along with its famous 
three dark lanes, which give it its name and which radiate outwards from 
the central object: an 08-type star, which is the power source for the nebula. 
The northern nebulosity is in fact a reflection nebula, and so it is harder to 
observe. 

Messier 8 NGC 6523 18 h 03.8 m -24° 23' May-Jun-Jul 

IS 1-5 ©45^430' Sagittarius 


50 


Astrophysics is Easy 


Also known as the Lagoon Nebula, this is the premier emission nebula of the 
summer sky. Binoculars will show a vast expanse of glowing green-blue gas 
split by a very prominent dark lane, whereas telescopes of aperture 30 cm will 
show much more intricate and delicate detail, including many dark bands. The 
Lagoon Nebula is located in the Sagittarius-Carina Spiral Arm of our Galaxy, at 
a distance of 5400 l.y. 

Messier 17 NGC 6618 18 h 20.8 m -16°11' May-Jun-Jul 

#1-5 ©20±416' Sagittarius 

Also known as the Swan or Omega Nebula. This is a magnificent object to 
see through binoculars, and perhaps a rival to the Orion Nebula, M42, for 
the summer sky. Not often observed by amateurs, as it is low down in the 
sky from, say, UK latitudes. With telescopes, the detail of the nebula becomes 
apparent, and with the addition of a light filter, it can in some instances 
surpass M42. Certainly, it has many more dark and light patches than its winter 
cousin, although it definitely requires an OIII filter for the regions to be fully 
appreciated. 

Messier 16 IC 4703 18 h 18.6 m -13°58' May-Jun-Jul 

#1-5 ©35±430' Serpens Cauda 

Also known as the Star Queen or Eagle Nebula, a famous though not often 
observed nebula. As usual, the use of a filter enhances its visibility. The “Black 
Pillar” and associated nebulosity are difficult to see, even though they are 
portrayed in many beautiful photographs (a prime example of astronomical 
imagery fooling the amateur into thinking that these justifiably impressive objects 
can easily be seen through a telescope). Nevertheless, it can be spotted by an 
astute observer under near-perfect conditions. 

Caldwell 27 NGC 6888 20 h 12.0 m +38°21' Jun-Jul-Aug 

#1-5 ©18' ±413' Cygnus 

Also known as the Crescent Nebula, this difficult nebula is included here as it is 
a prime example of several relevant phenomena associated with star formation. 
With good conditions, the emission nebula will live up to its name, having an 
oval shape with a gap in the ring on its southeastern side. The nebula is known as 
a Stellar Wind Bubble, and it is the result of a fast-moving stellar wind from a 
Wolf-Rayet star that is sweeping up all the material previously ejected during its 
red giant stage. 

IC 5067-70 20 h 50.8 m +44°21' Jul-Aug-Sep 

#1-5 ©25±410' Cygnus 

Also known as the Pelican Nebula, this nebula, close to the North American 
Nebula (see the entry below), has been reported to be visible to the naked eye 
and is easily glimpsed in binoculars as a triangular, faint, and hazy patch of 
light. Remember, it can be seen best with averted vision and the use of light 
filters. 

Caldwell 20 NGC 7000 20 h 58.8 m +44°12' Jul-Aug-Sep 

#1-5 ©120 ±4 100' Cygnus 


The Interstellar Medium 


51 


Also known as the North America Nebula. It is located just west of Deneb 
and is a magnificent site in binoculars, melding, as it does, into the stunning 
star fields of Cygnus. Provided you know where and what to look for, the 
nebula is visible to the naked eye. The dark nebula lying between it and the 
Pelican Nebula is responsible for their characteristic shape. Until recently, Deneb 
was thought to be the star responsible for providing the energy to make the 
nebula glow, but recent research points to several unseen stars’ being the power 
sources. 


IC 1396 21 h 39.1 m +57°30' Jul-Aug-Sep 

#3-5 ©170^40' Cepheus/ Cygnus 

One of the few emission nebulae visible to the naked eye (under perfect seeing 
conditions, of course!) and easily spotted in binoculars, it is an enormous patch 
of nebulosity, over 3°, spreading south of the orange star Mu (/jl) Cephei. As 
usual, a telescope will lessen the impact of the nebula but the use of filters will 
help to locate knots, patches of brighter nebulosity, and dark dust lanes. The use 
of dark adaption and averted vision will enhance the observation of this giant 
emission nebula. 


Caldwell 19 IC 5146 21 h 53.4 m +47° 16' Jul-Aug-Sep 

#3-5 ©12|12' Cygnus 

Also known as the Cocoon Nebula, Caldwell 19 has a low surface brightness and 
appears as nothing more than a hazy amorphous glow surrounding two 9th- 
magnitude stars. The dark nebula Barnard 168 (at the end of which the Cocoon 
lies) is surprisingly easy to find and can thus act as a pointer to the more elusive 
emission nebula. The whole area is a vast stellar nursery, and recent infrared 
research indicates the presence of many new protostars. 

Caldwell 11 NGC 7635 23 h 20.7 m +61° 12' Aug-Sep-Oct 

#1-5 ©16^9' Cassiopeia 

Also known as the Bubble Nebula, this is a very faint nebula, even in telescopes of 
aperture 20 cm. An 8th-magnitude star within the emission nebula and a nearby 
7th-magnitude star hinder its detection due to their combined glare. Research 
suggests that a strong stellar wind from a star pushes the material out (creating 
the “Bubble”) and heats up a nearby Molecular Cloud, which in turn ionizes the 
“Bubble.” 


NGC 604 01 h 33.9 m +30°39' Sep-Oct-Nov 

#3-5 ©60^35' Triangulum 

Now for something very special: the brightest emission nebula that can be 
glimpsed. It is actually in another galaxy — M33 in Triangulum. It appears as a 
faint hazy glow some 10' northeast of M33’s core. Oddly enough, due to M33’s low 
surface brightness (which often makes it a difficult object to find), the emission 


52 Astrophysics is Easy 

nebula may be visible while the galaxy is not! It is estimated to be about 1000 
times bigger than the Orion Nebula. 


Caldwell 49 NGC 2237-3906 h 32.3 m +05°03' Nov-Dec-Jan 

#1-5 ©80^»60' Monoceros 

Also known as the Rosette Nebula, this is a giant emission nebula that has the 
dubious reputation of being very difficult to observe. But this is wrong — on clear 
nights it can be seen with binoculars. Over 1° in diameter, it covers an area of 
sky four times larger than a full Moon! With a large aperture and light filters, 
the complexity of the nebula becomes readily apparent, and under perfect seeing 
conditions dark dust lanes can be glimpsed. The brightest parts of the emission 
nebula have their own NGC numbers: 2237, 2238, 2239, and 2246. It is a young 
nebula, perhaps only half a million years old, and star formation may still be 
occurring within it. Photographs show that the central area contains the star 
cluster NGC 2244, along with the “empty” cavity caused by the hot young stars 
blowing the dust and gas away. Also known as the Rosette Molecular Complex 
(RMC). 


NGC 2024 05 h 40.7 m 02°27' Nov-Dec-Jan 

#2-5 ©30^30 Orion 

This nebula lies next to the famous star Zeta Orionis, which is unfortunate as 
the glare from the star makes observation difficult. It can, however, be glimpsed 
in binoculars as an unevenly shaped hazy and faint patch to the east of the star, 
provided the star is placed out of the field of view. With large telescopes and 
filters, the emission nebula is a striking object and has a shape reminiscent of a 
maple leaf. 


Caldwell 46 NGC 2261 06 h 39.2 m +08° 44' Nov-Dec-Jan 

#1-5 ©3. 5^1. 5' Monoceros 

Also known as Hubble’s Variable Nebula, Caldwell 46 is easily seen in telescopes 
of 10 cm as a small, comet-like nebula, which can be seen even from the suburbs. 
Larger apertures just amplify what is seen with little detailed visibility. What 
we see is the result of a very young and hot star clearing away the debris 
from which it was formed. The star R Monocerotis (buried within the nebula 
and thus invisible to us) emits material from its polar regions, and we see the 
north polar emissions, with the southern emission blocked from view by an 
accretion disc. The variability of the nebula, reported in 1916 by Edwin Hubble, 
is due to a shadowing effect caused by clouds of dust drifting near the stars. 
It was also the first object to be officially photographed with the 200-inch Hale 
Telescope. 


NGC 1554-55 04 h 21.8 m +19°32' 

2-5 ©1^7 variable' 


Oct— Nov— Dec 
Taurus 


The Interstellar Medium 


53 


Also known as Hind’s Variable Nebula. I decided to include this object despite 
its difficulty to locate and observe because it is so interesting. This famous but 
incredibly faint emission nebula is located to the west of the famous star T Tauri, 
the prototype for a class of variable star. The nebula was much brighter in the 
past, but it is now an exceedingly difficult object to locate. With a large aperture, 
it will appear as a small faint hazy patch. When (and if!) located, it does bear 
higher magnification well. It may become brighter in the future, so it is worth 
looking for in the hope that it makes a reappearance. 


2.4 Dark Nebulae 


Dark nebulae (also known as dense clouds, briefly mentioned earlier in the 
chapter) by nature differ from other nebulae in one major respect: they do not 
shine. In fact, when you observe them, you are actually not seeing them by any 
light-emitting process, but rather for their light-blocking ability. They are vast 
clouds of gas molecules, such as H 2 , HCN, OH, CO, and CS, as well as dust grains. 
These grains, however, bear no resemblance to the dust we see on Earth. They 
are microscopic in size, believed to be in the region of 20 to 30 nm. However, 
ice (either water ice — H 2 0 — or ammonia ice — NH 3 ) may condense on them, 
forming a “mantle,” which then increases their size up to 300 nm. Dust grains are 
shaped like long spindles, and in some cases, they rotate. The actual composition 
of the grains is a topic of vigorous debate, but they are believed to be made, 
in various unknown amounts, from carbon in the form of graphite, along with 
silicon carbides, and silicates of magnesium and aluminium. 

The formation of the dust grains is spectacular! They are believed to have 
been formed in the outer regions of stars — in particular, the cool supergiants, 
and the R Corona Borealis-type stars. Dense molecular clouds are also a possible 
formation site. The temperature of the grains is thought to be about 10 to 100 K, 
which is cool enough to allow the formation of molecules. In a typical dark 
nebula, there may be anywhere from 10 4 to 10 9 particles made up of atoms, 
various molecules, and dust grains. 

Due to their vast size, the nebulae appear dark and so are very effective at 
scattering all the light, with the result that hardly anything reaches the naked 
eye. The process of scattering the light is so effective that, for instance, visible 
light emitted from the center of our Galaxy is nearly 100% extinguished by the 
dust clouds between us and its center. This is why the appearance of the central 
region in visible light is still a mystery. The scattering and absorption of light is 
known as extinction. Do not be confused by thinking that these clouds of dust 
grains are very dense objects. They are not. Most of the material in the cloud is 
molecular hydrogen (along with carbon monoxide, which is responsible for their 
radio emission), and the resulting density is low. There is also some evidence 
to suggest that the dust grains present in the clouds have different properties to 
those in the interstellar medium. 

Many dark nebulae are actually interacting with their environs, as witnessed by 
the spectacular images taken by the Hubble Space Telescope of M16 in Serpens. 




54 


Astrophysics is Easy 


The images show dust clouds containing dense regions, or globules, resisting 
the radiation pressure from close, hot young stars, with the result that many 
of the globules are trailing long tails of material. The area near the Horsehead 
Nebula in Orion is also famous for its image of the radiation from the super- 
giant stars of Orion’s Belt impacting on the dark clouds to either side of the 
Horsehead, with the result that material is ionized and streaming from the cloud’s 
surface. 

Most dark clouds have vastly different shapes, and this is for several reasons. 
It may be that the cloud was originally spherical in shape, and thus it would 
have presented a circular image to us, but hot stars in its environment disrupted 
this by radiation pressure and stellar winds. Shock fronts from nearby super- 
novae can also have an impact. The gravitational effects from other clouds, 
stars, and even that of the Milky Way itself all have a role to play in deter- 
mining the shape of a cloud. It is also believed that magnetic fields may have 
some limited effect. As many of these dark clouds are part of a much larger 
star-forming region, the new stars will themselves influence and alter their 
shape. 

Let us now look at a few examples of dark nebulae. The “opacity” of a dark 
nebulae is a measure of how opaque the cloud is to light, and thus how dark it 
will appear. There is a rough classification system that can be used; a value of 
1 for a dark nebula indicates that it very slightly attenuates the starlight from 
the background Milky Way; conversely, a value of 6 means that the cloud is 
nearly black, and it is given in the symbol ♦ . Observing dark nebulae can be a 
very frustrating pastime. The best advice I can offer is to always use the lowest 
possible magnification. This will enhance the contrast between the dark nebula 
and the background star field. If a high magnification is used, the contrast will be 
lost, and you will only see the area surrounding the dark nebula, not the nebula 
itself. Dark skies are a must with these objects, as even a hint of light pollution 
makes their detection an impossible task. 


This is a long band of dark nebula, easily spotted in binoculars as lying halfway 
between Psi (v|r) and Chi (y) Lupi. It is best seen in low-power, large-aperture 
binoculars, as it stands out well against the rich background of star field. 


2.4.1 Famous Dark Nebulae 


Barnard 228 
+ 6 


©240^20' 


15 h 45.5 m —34° 24' Apr-May-Jun 

Lupus 


Barnard 59, 65-7 LDN 1773 17 h 21.0 m 


— 27°23' May-Jun-Jul 
Ophiuchus 


♦ 6 


©300^60' 


Also known as the Pipe Nebula (Stem) and Lynds Dark Nebula 1773, this large 
dark nebula is visible to the naked eye because it stands out against a star-studded 


The Interstellar Medium 


55 


field and is best viewed with lower-power binoculars. With the unaided eye, it 
appears as a straight line, but under magnification its many variations can be 
glimpsed. 


Barnard 78 LDN 42 17 h 33.0 m -26°30' May-Jun-Jul 

+ 5 ©200^+150' Ophiuchus 

Also known as the Pipe Nebula (Bowl). This is part of the same dark nebula 
as above. The bowl appears as a jagged formation, covering over 9°. The whole 
region is studded with dark nebulae and is thought to be part of the same 
complex as that which encompasses Rho (p) Ophiuchi and Antares, which are 
more than 700 l.y. away from it. 


Barnard 86 LDN 93 18 h 03.0 m -27°53' May-Jun-Jul 

♦ 5 ©6' Sagittarius 

Also known as the Ink Spot, Barnard 86 is located within the Great Sagittarius Star 
Cloud. It is a near-perfect example of a dark nebula, appearing as a completely 
opaque blot against the background stars. 

Barnard 87, 65-7 LDN 1771 18 h 04.3 m -32°30' May-Jun-Jul 

4-4 ©12' Sagittarius/ Ophiuchus 

Also known as the Parrot Nebula, this is not a distinct nebula, but it stands 
out because of its location within a stunning background of stars. Visible in 
binoculars as a small, circular dark patch, it is best seen in a small telescope of 
about 10 to 15 cm. 


Lynds 906 20 h 40.0 m +42°00' Jul-Aug-Sep 

♦ 5 © Cygnus 

Also known as the Northern Coalsack, this is probably the largest dark nebula 
of the northern sky. It is an immense region, easily visible on clear, moonless 
nights just south of Deneb and lying just at the northern boundary of the Great 
Rift, a collection of several dark nebulae that bisects the Milky Way. The Rift is 
part of a spiral arm of our Galaxy. 


Barnard 352 20 h 57.1 m +45° 54' Jul— Aug— Sep 

+ 5 ©20^10' Cygnus 

Visible in binoculars as a well-defined triangular dark nebula, this is part of 
the much more famous North American Nebula; this dark part is located to the 
north. 


Barnard 33 05 h 40.9 m -02°28' Nov-Dec-Jan 

4-4 ©6^4' Orion 


56 


Astrophysics is Easy 


Also known as the Horsehead Nebula. Often photographed, but rarely observed, 
this famous nebula is very difficult to see. It is a small dark nebula that is seen in 
silhouette against the dim glow of the emission nebula IC 434. Both are very faint 
and require perfect seeing conditions. Such is the elusiveness of this object that 
even telescopes of 40 cm are not guaranteed a view. Dark adaption and averted 
vision, along with the judicious use of filters, may result in its detection, so 
have a go! 


2.5 Reflection Nebulae 


The final classification of nebulae is reflection nebulae. As the name suggests, 
these nebulae shine by the lights reflected from the stars within them, or from 
nearby stars. Like the emission nebulae, these vast clouds consist of both gas 
and dust, but in this case, the concentration of dust is far less than that found 
in emission nebulae. One of the characteristics of particles, or grains that are 
so small (in proportion to the wavelength of light), is their property of selec- 
tively scattering light of a particular wavelength. If a beam of white light shines 
upon a cloud containing the grains, the blue light is scattered in all directions, 
a phenomenon similar to that seen in the Earth’s sky 11 (hence its blue color). 
This is one reason that reflection nebulae appear so blue in photographs; it is 
just the blue wavelengths of the light from (usually) hot blue stars nearby. To 
be scientifically accurate, the nebulae should be called scattering nebulae instead 
of reflection nebulae, but the name has stuck. An interesting property of the 
scattered light is that the scattering process itself polarizes the light, which is 
useful in the studies of grain composition and structure. 

But that’s not all... if a star that lies behind a dust cloud is observed, some of 
its blue light is removed by the process discussed above, and an effect known as 
interstellar reddening occurs, which makes the light from the star appear redder 
than it actually is. This leads to a further phenomenon associated with dust 
grains called interstellar extinction, which should be mentioned because it affects 
all observations. Astronomers noticed that the light from distant star clusters 
was fainter than expected, and this was due to dust lying between us and the 
cluster. This in fact makes all objects fainter than they actually are and leads to 
an underestimation of their luminosity and an overestimation of their distance. 
Thus, interstellar extinction must be taken into account when making detained 
measurements. 

Several reflection nebulae reside within the same gas clouds as emission 
nebulae. The Trifid nebula is a perfect example. The inner parts of the nebula 
are glowing with a telltale pink color, indicative of the ionization process respon- 
sible for the emission, whereas further out from the center, the edge material is 
definitely blue, signposting the scattering nature of the nebula. Visually, reflection 
nebulae are very faint objects with a low surface brightness, so they are not easy 
targets. Most require large-aperture telescopes with moderate magnification to be 
seen, but a few are visible in binoculars and small telescopes. Note that excellent 
seeing conditions and very dark skies are required. 



The Interstellar Medium 


57 


2.5.1 Brightest Reflection Nebulae 

NGC 1435 03 h 46.1 m +23°47' Oct-Nov-Dec 

#2-5 ©30^30' Taurus 

Also known as Tempel’s Nebula, this faint patch of reflection nebula is located 
within the most famous star cluster in the sky — the Pleiades. The nebula itself 
surrounds the star Merope, one of the brighter members of the cluster, and under 
perfect conditions can be glimpsed with binoculars. Several other members of 
the cluster are also enshrouded by nebulosity, but these require exceptionally 
clear nights and, incidentally, clean optics, as even the slightest smear on, say, 
a pair of binoculars will reduce the chances to nil. 


Caldwell 31 IC 405 05 h 16.2 m +34°16' Nov-Dec-Jan 

#2-5 ©30^19' Auriga 

Also known as the Flaming Star Nebula, this is a very challenging reflection 
nebula to observe. It is actually several nebulae, including IC 405, 410, and 417, 
plus the variable star AE Aurigae. The use of narrow-band filters will be justified 
with this reflection nebula, as they will highlight the various components. 


2.6 Molecular Clouds 


We have seen that interstellar space is filled with gas and dust, and that in certain 
locations, concentrations of this material gives rise to nebulae. But the location 
of these nebulae are not, as one might expect, entirely random. The areas that 
give rise to star formation are called Molecular Clouds. These clouds are cold, 
perhaps only a few degrees above absolute zero, and occupy enormous regions of 
space. Due to the conditions within them, molecular clouds allow the formation 
of several molecules [e.g., carbon monoxide (CO), water (H 2 0), and hydrogen 
molecules (H 2 ) 12 ]. Although the most abundant molecule in a cloud, molecular 
hydrogen is very difficult to observe because of the low temperature. On the other 
hand, CO can be detected when certain portions of the cloud are 10-30 K above 
absolute zero. It is these molecules that allowed molecular clouds to be discovered 
by two radio astronomers — Philip Solomons and Nicholas Scoville — who, in 
1974, found traces of the carbon monoxide molecule in the Galaxy. 

Molecular clouds are truly gigantic and contain vast amounts of hydrogen. 
They can have masses from 10 5 to 2 x 10 6 solar masses, and diameters anywhere 
from 12 to 120 pc, or about 40 to 350 l.y. The total mass of molecular clouds in 
our Galaxy is thought to be about 5 billion solar masses. 13 But even though these 
molecular clouds are so vast, do not be fooled into thinking that we are talking 
about something that resembles, in structure, conditions similar to a foggy day, 
with hydrogen and dust being so dense that you hardly see anything in front 
of you. If we could go inside one of these clouds, there would be about 200 or 
300 hydrogen molecules per cubic centimeter. This is not a lot, even though it is 




58 


Astrophysics is Easy 


several thousand times greater than the average density of matter in our Galaxy. 
Even more staggering, it is 10 17 times less dense than the air we breathe. 

Astronomers have deduced that molecular clouds and CO emission are 
intimately linked, and by looking at the areas in our Galaxy where CO emission 
originates, we are in fact looking at those areas where star formation is taking 
place. Because the molecular clouds are, by comparison with the rest of the ISM, 
heavy and dense, they tend to settle toward the central layers of the Milky Way. 
This has produced a phenomenon we have all seen — the dark bands running 
through the Milky Way. Surprisingly, it was found that the molecular clouds 
in which star formation occurs outline the spiral arms of the Galaxy and lie 
about 1000 pc apart, strung out along the arms rather like pearls on a necklace. 14 
However, spiral arms of galaxies are not the only place where star formation can 
occur. There are several other mechanisms that can give rise to stars, as we shall 
see in the chapter on stars. 


2.7 Protostars 


I have included the topic of protostars in this chapter and not the following 
because we are still discussing large, diffuse clouds of gas and dust, albeit briefly, 
before they are turned into proper stars. So let’s begin by looking at the mecha- 
nisms by which stars are believed to have been formed. 

We have discussed the fact that space is full of gas and dust, and that local 
concentrations of this material give rise to nebulae. But how do stars form in 
these regions? It may seem obvious in hindsight that a star will form in those 
clouds where the gas and dust are particularly dense and thus will allow gravity 
to attract the particles. An additional factor that will assist in formation is a very 
low temperature of the cloud. A cold cloud means that the (thermal) pressure of 
the ISM is low. A cold temperature is not only helpful but in fact a prerequisite, as 
clouds have a high (thermal) pressure, which tends to overcome any gravitational 
collapse. It is a delicate balancing act between gravity and pressure, whereby if 
gravity dominates, stars form. 

From our earlier discussion in the book, you should have realized by now that 
there is only one place where conditions like those just mentioned arise: the dark 
nebulae. As the cloud contracts, pressure and gravity permitting, the dust and gas 
cloud becomes very opaque and the precursor region to star formation. These 
regions are often called Barnard objects, after the astronomer who first catalogued 
them, Edward Barnard. 15 There are also even smaller objects, sometimes located 
within a Barnard object. These resemble small, spherical dark blobs of matter 
and are referred to as Bok globules, named after astronomer Bart Bok. It may 
help you to think of a Bok globule as a Barnard object but with its outer layers, 
which are the less-dense regions, dispersed. 

Radio measurements of Bok globules indicate that their internal temperature 
is a very low 10 K, and their density, although only about 100 to 20,000 particles 
(dust grains, gas atoms, and molecules) per cubic centimeter, is considerably 
greater than that found in the ISM. The size of these objects can vary considerably; 
there are no standard sizes, but on average, a Bok globule is about 1 pc in 




The Interstellar Medium 


59 


diameter, ranging anywhere from 1 to 1000 M 0 . The larger Barnard object, on the 
other hand, can have a mass of about 10, 000 M 0 , with a diameter of about 10 pc. 
As you can imagine, the sizes of these objects vary greatly and are determined 
by the local conditions in the ISM. 

Now, if conditions permit, the densest areas within these objects and globules 
will further contract under gravitational attraction. A consequence of this 
contraction is the heating up of the blob material; however, the cloud can 
radiate this thermal energy away, and in doing so prevent the pressure from 
building up enough to resist the contraction. During the early phase of collapse, 
the temperature remains below 100 K, and the thermal energy is transported 
from the warmer interior to the exterior of the cloud by convection, causing 
the cloud to glow in infrared radiation. This ongoing collapse has the effect of 
increasing the cloud’s density, but this makes it difficult for the radiation to 
escape from the object. Consequently, the central regions of the cloud become 
opaque, which traps nearly all the thermal energy produced by the gravita- 
tional collapse. Trapping the energy results in a dramatic increase in both 
pressure and temperature. The ever-increasing pressure fights back against the 
overpowering crush of gravity, and the now-denser fragment of cloud becomes 
a protostar — the seed from which a star is born. At this stage, a protostar may 
resemble a star, but it is not really a star, as no nuclear reactions occur in 
its core. 

The time taken for the above scenario to occur can be extremely short, in an 
astronomical sense — maybe of the order of a few thousand years. The protostar is 
still quite large. For example, after, say, 1000 years, a protostar of 1 M Q can be 20 
times larger than the Sun’s radius, R 0 , and about 100 times as luminous, 100 L 0 . 


2.8 The Jeans Criterion 


You might think from the previous sections that star formation is a pretty 
straightforward process, and that if there is enough material (i.e., gas and dust) 
and a long enough period of time, the only possible outcome is the formation of 
a star. You would be wrong! 

Remember that an interstellar cloud, however large (or small), performs a 
delicate balancing act between the gravitational attraction from all the cloud’s 
particles, which is trying to collapse the cloud, and the thermal energy (think 
of it as the cloud’s heat), which is trying to resist this collapse. If one is more 
dominant than the other, a star may form. 

The question to ask yourself is, “What decides whether gravity wins?” 
This is where the Jeans Criterion 16 comes into play. In a cloud with a 
specific density, temperature, and mass, these criteria describe the small- 
sized cloud and its minimum mass where gravity could overcome the 
thermal pressure and so result in collapse. As you can imagine, some 
quite involved equations are used; however, we can make approximations 
of them (see Box 2.1). 




60 


Astrophysics is Easy 


The critical mass of the cloud is known as the Jeans Mass, M,, and the critical 
size, the Jeans Length R y The Jeans Mass is the mass of the cloud whose radius 
is the Jeans Length. 

From the above description, you can see that there are a few conditions that 
make cloud collapse more likely: the cloud has a very low temperature (the 
cooler the cloud, the better the chances of collapse), and the cloud is more dense 
(a dense cloud has a better chance of collapse than one that is very thinly spread 
out). Thus, the dark, dense clouds discussed earlier would be ideal locations for 
cloud collapse. Indeed, in the darkest, densest clouds, only a few solar masses of 
material are necessary for collapse. 

Some dense, dark clouds have within them even denser areas, called clumps and 
cores, which may have masses ranging from 0.3 to 10 3 solar masses, and thus can 
satisfy the Jeans Criterion on their own. So now we have the situation of a large, 
dense, dark cloud’s collapsing, while inside it, there are clumps collapsing, as well! 

But, as you can imagine, things are far more complicated than the picture 
I have just drawn for you. As clouds, cores, and clumps collapse, they tend to 
warm up. This acts to inhibit the gravitational collapse; however, this brief hiccup 
is overcome, and collapse continues. 

One point that needs to be mentioned is that most of the diffuse clouds in the 
ISM are not close to the Jeans Criterion, so some sort of mechanism, or trigger, 
is needed to change the conditions. In fact, what is needed is something that will 
increase a cloud’s density (i.e., an event that will compress the cloud material 
into a smaller volume of space). 17 Once a trigger pushes the cloud closer to, and 
possibly over, the Jeans limits, then cloud collapse can begin (we discuss these 
possible triggers later in the book). 

Finally, imagine a massive cloud that does not initially satisfy the Jeans Criterion, 
but then something causes the cloud to collapse. Areas within the large cloud may 
now satisfy the criteria and so they themselves start to contract. In a cloud of several 
hundred to several thousand solar masses, there can be a lot of clumps, and this 
fragmentation, as it is called, could eventually give rise to a cluster of stars. Thus, 
this maybe a possible scenario for the formation of open star clusters. 

The Jeans Criterion is a good starting point in the description of cloud 
collapse, and today there exist far more sophisticated models that perhaps more 
accurately describe what is going on. Nevertheless, as a starting point, they 
adequately describe the possible beginning of star formation (see Box 2.1). 

Let us now move on to those objects we (hopefully) can see each and every 
night — stars. 


Box 2.1 : The Jeans Length and Jeans Mass 

The Jeans Length is approximately given by: 

R,k (kT/Gm 2 n) 1/2 

k is Boltzmann constant = 1.3806 x 10 -23 JK' 1 
T is temperature in K 



The Interstellar Medium 


61 


G is gravitational constant = 6.67 x 10 -11 Nm 2 kg~ 2 
m is mass of hydrogen atom = 1.67 x 10 _27 kg 
n is number of particles (number density) 

Example: 

If an interstellar cloud has a temperature of 50 K, and there are 10 11 hydrogen atoms 
per cubic meter, determine the Jeans Length and Jeans Mass. 

Using the above formula, we get: 

[ (1.38 x 10“ 23 ) x (50) ~| 1/2 

' ~ (6.67 x lO' 11 ) x (1.67 x 10~ 27 ) 2 x (10 11 ) 

R r ^6 x 10 15 m 
^0.2 pc 

The Jeans Mass can easily be estimated by multiplying the density by the volume: 

Mj = (4t7/3)(1.67 x 10~ 27 )(10 u )(6 x 10 15 ) 3 
1.5 x 10 32 kg 
^76M 0 

Thus, in a cloud with a temperature of 50 K that has 10 11 atoms per cubic meter, 76 
solar masses of material is the minimum amount needed for gravitation to overcome 
any thermal pressure, with a radius of about 0.2 pc. 


Notes 


1. An important distinction is that n is not the same as the number of hydrogen 
atoms per cubic meter. If the hydrogen is in a molecular form, H 2 , then 
the number of separate particles is n H / 2. 

2. HII is pronounced "aitch 2.” 

3. Recall that the ISM is made up of about 74% hydrogen (by mass), 25% 
helium, and the rest metals. 

4. And in some cases, star death, namely supernova remnants, covered later. 

5. Our simple model of an atom has a central nucleus with electrons orbiting 
around it, somewhat like planets orbiting a Sun. Electrons with a lot of 
energy are in the outer orbits, while electrons with less energy are closer 
to the nucleus. Not all orbits are allowed by quantum mechanics: to move 
up to higher energy levels, electrons need a very specific amount of energy; 
too much or too little, and an electron will not move. 

6. The time spent before recombining is very short — millionths of seconds — 
but also depends on the amount of radiation present and the density of the 
gas cloud. 

7. Unfortunately, the red glow is usually too weak to be seen at the eyepiece. 




62 Astrophysics is Easy 

8. These lines are a rich blue-green color and, under good seeing conditions 
and with clean optics, can be glimpsed in the Orion Nebula, M42. 

9. In some astrophysical contexts, such as in the center of quasars, conditions 
exist that can give rise to terms such as Fe23. The amount of radiation is so 
phenomenal that the atom of iron (Fe) has been ionized to such an extent, 
it has lost 22 of its electrons! 

10. This is often called the Stromgren sphere, named after the astronomer 
Bengt Stromgren, who did some pioneering work on HII regions. 

11. Note that scattering of water molecules, and not dust, is responsible for the 
blue sky on Earth. 

12. Other molecules such as ammonia (NH 3 ) and alcohol (CH 3 OH!!!) have also 
been detected. 

13. In areas where the average density exceeds, say, a million solar masses, 
clouds referred to as Giant Molecular Clouds can form. 

14. Molecular clouds can be found outside of spiral arms, but current ideas 
suggest that the spiral arms are regions where matter is concentrated, due 
to gravitational forces. The molecular clouds pass through the arms and 
are “squeezed.” This dense region then gives rise to star-forming regions. 

15. See the section on dark nebulae for observable examples of Barnard objects. 

16. James Jeans (1877-1946) was a British astronomer, and the first person to 
mathematically describe the necessary conditions for star collapse. 

17. At this point, the increase in density is thought to be a more important 
condition than a commensurate increase in temperature. 


CHAPTER THREE 


Stars 


3.1 The Birth of a Star 


A newly born star can be thought of as having been born when the core 
temperature of the protostar reaches about 10 million K. At this temperature, 
hydrogen fusion can occur efficiently by the proton-proton chain. 1 The moment 
the ignition fusion process occurs will halt any further gravitational collapse of 
the protostar. The star’s interior structure stabilizes, with the thermal energy 
created by nuclear fusion maintaining a balance between gravity and pressure. 
This important act of balancing is called gravitational equilibrium. 2 It is also 
sometimes referred to as hydrostatic equilibrium. The star is now a hydrogen- 
burning main-sequence star. 

The time taken for the formation of a protostar to the birth of a main-sequence 
star depends on the star’s mass. This is an important point to emphasize. A star’s 
mass determines a lot! A handy reference to remember is that massive stars do 
everything faster! A high-mass protostar may collapse in only a million years or 
less, while a star with a mass of 1 M 0 could take around 50 million years. A star 
with a very small mass, say, an M-type star, could take well over 100 million 
years to collapse. This means that very massive stars in a young star cluster may 
be born, live, and die before the very smallest stars finish their infant years! 

The changes, or transitions, that occur to a protostar’s luminosity and surface 
temperature can be shown on a special H-R diagram. This is known as an 
evolutionary track, or a lifetrack, of a star. 3 Each point along the star’s track 
represents its luminosity and temperature at some point during its life, and so it 


63 


64 


Astrophysics is Easy 


shows us how the protostar’s appearance changes due to changes in its interior. 
Figure 3. 1 shows the evolutionary tracks for several protostars of different masses, 
from 0.5 M 0 to 15 M 0 (it is important to realize that these evolutionary tracks are 
theoretical models, and the predictions are only as good as theory; 4 they seem to 
work very well and are being improved all the time). Recall that protostars are 
relatively cool, and so the tracks all begin at the right side of the H-R diagram. 
However, subsequent evolution is very different for stars of differing mass. 

The evolutionary lifetracks of seven protostars are shown. Also identified are 
the stages reached after an indicated number of years of evolution (dashed 
lines). The mass shown for each protostar is the final mass when it becomes a 
main-sequence star. Note that the greater the mass, the higher the temperature 
and luminosity. 

As an example of an evolutionary lifetrack for a protostar, we shall look at the 
lifetrack for a 1M 0 , rather like the Sun. This period in the star’s life has 4 very 
distinct phases: 



Figure 3.1. Pre-main-sequence lifetracks. 



Stars 


65 


Phase 1: The protostar forms from a cloud of cold gas and thus is on the far 
right side of the H-R diagram; however, its surface area is enormous, 
and so its luminosity can be very large — it may be 100 times more than 
luminous when it becomes a star. 

Phase 2: Due to its large luminosity, the young protostar rapidly loses the energy 
it generated via gravitational collapse, and so further collapse proceeds 
at a relatively rapid rate. Its surface temperature increases slightly during 
the next several million years, but its diminishing size reduces the 
luminosity. The evolutionary track now progresses almost vertically 
downward on the H-R diagram. 

Phase 3: Now that the core temperature has reached 10 million K, hydrogen 
nuclei fuse into helium. The rate of nuclear fusion, however, is not 
sufficient to halt the collapse of the star, although it is slowed down 
considerably. As the star shrinks, its surface temperature increases. 
The process of shrinking and heating will result in a small increase in 
luminosity over the next 10 million years. The evolutionary track now 
progresses leftward and slightly upward on the H-R diagram. 

Phase 4: Both the rate of nuclear fusion and the core temperature increase 
over the next tens of millions of years. Once the rate of fusion is 
high enough, gravitational equilibrium is achieved, and fusion becomes 
self-sustaining. The result is that the star settles onto the hydrogen- 
burning main sequence (see fig 3.2). 


From the viewpoint of an observer, this stage of stellar evolution does not 
present itself with many visible objects. Even though the luminosity of such 


Luminosity (Lq) 


K 


] 




Phase 1 : A cloud of gas and dust become a protostar 
when radiation no longer escapes from the interior of the 
cloud 


40,000 20,000 10,000 5000 

< — Surface temperature (K) 


5000 


3000 


Figure 3.2. The evolutionary track of a protostar. 



66 


Astrophysics is Easy 


objects is very high, we will never see one. The reason is obvious: they are 
enshrouded within vast clouds of interstellar dust, which, if you recall, are very 
efficient at blocking out light. The dust in the vicinity of a protostar, often called 
a cocoon nebula, absorbs the light, and so it is very difficult to observe at visible 
wavelengths. 5 On the other hand, they can be seen at infrared wavelengths... but 
this does not really help us, as visual observational astronomers on Earth. 


3.2 Pre-Main-Sequence Evolution 
and the Effect of Mass 


The previous sections explained how a cloud can contract and become a protostar. 
In fact, due to the immense amount of material in a molecular cloud, it is believed 
that rather than an individual protostar being formed, several are formed as a 
star cluster. However, there is a small problem with this scenario; at the time of 
writing this book, there is no satisfactory explanation for protostars of differing 
masses actually forming within the same cloud. What are the processes that 
govern the clumping and fragmentation of the cloud into protostars of widely 
differing masses? Even though we cannot explain the process, we can at least 
observe its results. 

Let us begin this section by looking at how protostars of differing mass are 
believed to have been formed. We’ll begin with a star of 1 M Q , a star just like Sun. 
The outer layers of such a protostar are cool and opaque, 6 which means that energy 
released as radiation due to the shrinkage of the inner layers cannot reach the 
surface. Thus, the only way of moving this energy toward the surface layers must 
be by the less-efficient and slower method of convection. The result of this process 
is that the temperature remains more or less constant as the protostar shrinks, 
while the luminosity decreases because the radius decreases, 7 and the evolutionary 
track moves downward on the H-R diagram. This is depicted in Figure 3.1. 

I said previously that the surface temperature remains roughly constant during 
this phase, but conditions inside the protostar are far from unchanging. The 
internal temperature starts to increase during this time, and the interior becomes 
ionized. This reduces the opacity within the protostar and allows the transfer of 
energy by radiation in the interior regions, and by convection in the outer layers. 
This process is the one that is ongoing within the Sun today. The net result of 
these changes is that energy can escape much more easily from the protostar, and 
thus the luminosity increases. This increase in energy transport is represented by 
the evolutionary track’s bending upward (meaning higher luminosity) and to the 
left (higher temperature). After an interval of a few million years, the temperature 
within the protostar is high enough — 10 million K — for nuclear fusion to begin 
and, eventually, enough heat and associated internal pressure are created so as 
to balance the gravitational contraction of the star. We can say that at this point, 
hydrostatic equilibrium has been reached, and the protostar has reached the 
main sequence — it is now a main-sequence star. 

As to be expected, a more massive protostar will evolve in a different way. 
Protostars with a mass of about or greater than 4M 0 contract and heat up 


Stars 


67 


at a more rapid rate, and so the hydrogen-burning phase begins earlier. The 
net result is that the luminosity stabilizes at approximately its final value, but 
the surface temperature continues to increase as the protostar continues to 
shrink. The evolutionary track of such a high-mass protostar illustrates this on 
the H-R diagram; the luminosity is nearly horizontal (meaning nearly constant 
luminosity) from right to left (increasing surface temperature). This is especially 
true for the stars at mass 9M 0 and 15M 0 . 

An increase in mass will result in a corresponding increase in pressure and 
temperature in the interior of a star. This is very significant because it means that 
in very massive stars, there is a much greater temperature difference between the 
core and its outer layers as compared to, say, the Sun. This allows convection 
to occur much deeper into the star’s interior regions. In contrast, the massive 
star will have very low-density outer layers, and so energy flow in these regions 
is more easily performed by radiative methods than by convective methods. 
Thus, stars on the main sequence with a mass greater than about 4 M Q will have 
convective interiors and radiative outer layers, while stars less than about 4M G 
will have radiative interior regions and convective outer layers. 

At the very low end of the mass scale, stars that have a mass less than about 
0.8 M 0 have a very different internal structure. In these objects, the interior 
temperature of the protostar is insufficient to ionize the inner region, which 
is thus too opaque to allow energy transport by radiation. The only possible 
method to transport the energy to the outer layers is by convection. In these stars, 
convective methods are the only means of energy transport. Examples of the 
interior structures of low-mass, high-mass, and very low-mass stars are shown 
in Figure 3.3. 

Energy flows from the core by convection in the inner regions and by radiation 
in the outer layers in stars of mass greater than 4M 0 . 

Energy flows outward from the core by radiative means in inner regions and 
by convection in outer layers in stars with a mass of less than 4 M Q and greater 
than 0.8 M 0 . 

Energy flows outward by convection throughout the interior of the stars with 
a mass of less than 0.8 M 0 . 

A very important point to make here is that all the evolutionary tracks shown 
in Figure 3.1 end at the main sequence. Thus, the main sequence represents 
those stars in which nuclear fusion reactions are producing energy by converting 
hydrogen to helium. For the large majority of stars, this is a stable situation, and 
this endpoint on the main sequence can be represented by a Mass-Luminosity 
Relationship, which is shown in Figure 3.4. What this diagram implies is that the 
hot bright blue stars are the most massive, while the cool faint red stars are the 
least massive. 8 Thus, the H-R diagram is a progression not only in luminosity 
and temperature but in mass, as well. This can be succinctly summed up as “the 
greater the mass, the greater the luminosity.” 

For stars on the main sequence, there is a ratio between the mass and 
luminosity. Basically, the more massive the star, the greater its luminosity. A star 
of mass 10 M 0 has about 3000 L Q ; a star of mass 0.1 M Q has a luminosity of only 
0.001 L 0 . 

Now that we have discussed how stars are formed, and how star birth is 
depicted in the H-R diagram, it is important that I emphasize two factors that 


68 


Astrophysics is Easy 



Convection layers 


Radiative layers 



ii) Mass lying between 0.8 M s and 4 M s 


Convection layer 



iii) Mass less than 0.8 M s 


Figure 3.3. Mass of main-sequence stars. 



Figure 3.4. Mass-luminosity relationship. 




Stars 


69 


can cause confusion. First, if you look at the evolutionary track of protostars, 
several of them, especially the high-mass protostars, begin in the upper-right 
region. But they are not red giant stars ! The red stars are at a stage in their 
lives that occurs only after being a main-sequence star. The second point to note 
is that most stars spend most of their lives on the main sequence and only a 
relatively brief time as protostars. For example, a 1 M Q protostar takes about 
20 million years to become a main-sequence star, while a 12 M 0 may only take 
20,000 years. In contrast, a star like the Sun has been a main-sequence star for 
nearly 5 billion years and will remain one for another 5 billion! 

One final point is that the masses of stars have limits. Using theoretical models, 
astronomers have deduced that stars above 150-200 M 0 cannot form; they 
generate so much energy that gravity cannot contain their internal pressure. 
These stars literally tear themselves apart. At the other end of the scale, there 
is also, not surprisingly, a lower limit. Those stars with a mass of less than 
0.08 M 0 9 can never achieve the 10 million K core temperature necessary to 
initiate nuclear fusion. So, what is actually formed can be thought of as a 
“failed star,” which will slowly radiate away all of its internal energy, gradually 
cooling with time. These objects have been called Brown dwarfs and seem to 
occupy a strange area between what we think of as a planet and a star. Brown 
dwarfs radiate in infrared, making them very difficult to detect. The first known 
detection was in 1995 of Gliese 229B, a 0.05 M 0 object. Many astronomers 
believe these small, elusive objects are far more common than previously 
thought and may in fact be the most common form of ordinary matter 10 in the 
universe. 

From an observational point of view, this period in a star’s life does not present 
us with many observable opportunities. The protostars are cocooned within vast 
clouds of gas and dust and are therefore invisible to us. Some objects are, of 
course, visible if infrared telescopes are used. However, it is always worthwhile to 
look at areas of the night sky where we know such objects exist, even though they 
cannot be seen. We can always use our imagination as we gaze at them, knowing 
that hidden deep in these clouds are stars in the process of being formed. One 
such location is, of course, the Orion Nebula. 

Messier 42 NGC1976 05 h 35.4 m -05°27' Nov-Dec-Jan 

#1-5 ©65^60' Orion 

Also known as the Orion Nebula. This is the premier emission nebula and one 
of the most magnificent objects in the sky. It is part of the vast Orion complex, 
which contains star-forming regions, molecular clouds, and all sorts of nebulae! 
Visible to the naked eye as a barely resolved patch of light, it shows detail 
from the smallest aperture upwards. In binoculars, its pearly glow will show the 
structure in detail, and in telescopes of aperture 10 cm, the whole field will be 
filled. The entire nebulosity is glowing due to the light (and thus energy) provided 
by the famous Trapezium stars located within it. These stars are stellar power 
houses, pouring forth vast amounts of energy, and they are fairly new stars. What 
is also readily seen along with the glowing nebula are the dark, apparently empty 
and starless regions. These are still part of the huge complex of dust and gas but 
are not glowing by the process of fluorescence — instead, they are vast clouds of 
obscuring dust (the dark nebulae mentioned previously). The emission nebula is 


70 


Astrophysics is Easy 


one of the few that show definite color. Many observers report seeing a greenish 
glow, along with pale grey and blue, but to observe any color besides gray will 
require excellent observing conditions. Also, amateurs state that with very large 
apertures of 35 cm, a pinkish glow can be seen. Located within the nebula are the 
famous Kleinmann-Low Sources and the Becklin-Neugebauer Object, which are 
believed to be dust-enshrouded young stars. The whole nebula complex is a vast 
stellar nursery. M42 is at a distance of 1700 l.y. and about 40 l.y. in diameter. Try 
to spend a long time observing this object — you will benefit from it, and many 
observers just let the nebula drift into the field of view. 


3.3 Mass Loss and Gain 


3.3.1 TTauri Stars 

After reading the previous sections, you may have gotten the idea that star 
formation is simply a matter of material falling inward due to gravity. In fact, 
most of the material that makes up a cloud is ejected into space and never 
forms stars at all. This ejected material can help sweep away the gas and dust 
surrounding the young stars and make them visible to us. Several examples of 
such a process can be seen in the Rosette Nebula, the Triffid Nebula, and the 
Bubble Nebula, mentioned earlier. 

There are also examples of individual objects that eject material into space 
during this aspect of a star’s birth. These are called T Tauri stars, which are 
protostars whose luminosity can change irregularly in a matter of just a few 
days, and which also have both absorption and emission lines in their spectrum. 
In addition, due to the conflict between gravitational contraction and hydrogen- 
burning in these first stages of main-sequence stability, the element lithium 
is produced. Spectral lines of lithium are a signature of protostars of the T 
Tauri type. The masses of these stars are less than about 3M 0 and they seem 
to be about 1 million years old. If placed on an H-R diagram, they would 
be on the right-hand side of the main sequence. By analyzing the emission 
lines, we can see that surrounding these protostars are very thin clouds of 
very hot gas, which the protostar has ejected into space with a speed of about 
80km s _1 (300, 000km fT 1 ). A T Tauri star bears a superficial resemblance to the 
Sun in that it will exhibit a spectral type of F, G, or K, with a surface temperature of 
4000-8000 K. 

Over a period of a year, a typical T Tauri star would have ejected about 10 -8 to 
10“ 7 solar masses. You may think that this is a very small amount, but compared 
with the Sun, which loses about 10“ 4 M o a year, it is significant. This phase of a 
protostar, called the T Tauri phase, can last as long as 10 million years, during 
which it can eject roughly 1 M Q of material. A consequence of this is that the 
mass of the final main-sequence star is very much less than the mass it started 
with. As these are objects associated with star birth, they are often, if not always, 
found near or in the Milky Way. 

Other young stars with masses greater than 3 M Q do not vary in luminosity like 
T Tauri stars; they do eject mass, however, due to the extremely high radiation 


Stars 


71 


pressure at their surfaces. This class of star is called Ae or Be stars. Stars greater 
than 10 M q will reach the main sequence before the surrounding dust and gas 
from which they have been formed will have had a chance to disperse, and so 
these stars are often detected as highly luminous infrared objects located within 
the molecular clouds. 

Fortunately, as observers, there are several visible examples of T Tauri stars. They 
are, however, extremely faint, and so only the archetypal one is mentioned below. 

T Tauri 04 h 22.0 m +19°32' Oct-Nov-Dec 

8.5 13.6 v m dGe— Kle Taurus 

This star is about 1.8° west and slightlynorth of s (epsilon) Tauri, the northern- 
most bright star in the famous “v” shape of the Hyades star cluster. Discovered 
in 1852 by J. Hind (who also discovered the associated nebula — Hind’s Variable 
Nebula). The star varies irregularly in several aspects: the brightness varies from 
about 8th to 13th magnitude, the period, with a range from a few weeks to 
perhaps a few months, and the spectrum varying from G4 to G8. Oddly enough, 
the variation in spectral type does not necessarily correlate with variability in 
magnitude. T Tauri and the nebula lie within the Taurus Dark Cloud Complex, 
within which there are numerous but faint, variable nebulae and recently formed 
stars (other T Tauri and similar stars are VV Tauri and FU Orionis). 11 


3.3.2 Discs and Winds 

One aspect of protostar formation that came as a surprise to astronomers in 
the late twentieth century was a curious phenomenon observed in many young 
stars, including the T Tauri stars mentioned above. It involves a loss of mass, 
once again, but the loss is directed out from the young star in two jets; these 
are very narrow, usually flowing out along the rotation axis of the star and in 
opposite directions. This outflowing jet is referred to as a bipolar outflow. The 
material is moving with a velocity that can reach several hundred kilometers per 
second, and it sometimes interacts with the surrounding debris left over from 
star formation to form clumpy knots of material called Herbig-Haro objects. The 
lifetime of such a phenomenon is relatively short, maybe from 10,000 to 100,000 
years. The mechanism that forms these jets is not yet fully understood, although 
it is believed to involve magnetic fields. 

We have discussed mass loss in a protostar, but there exists a mechanism that 
can add mass to the normal star-formation process. Recall that a protostar is 
formed from falling gas and dust due to gravity. As clouds of denser material 
clump together, the protostar nebula will begin to rotate. This is a consequence 
of physics, and it is called the Conservation of Angular Momentum. The material 
will flatten itself out and form a disc, or protostellar disc, as it is called. The gas 
and dust particles within the nebula collide and spin inwards onto the forming 
protostar, thus adding to its mass. This process is often called accretion, and the 
build-up of material into the ever faster-rotating disc is called the circumstellar 
accretion disc. 


72 


Astrophysics is Easy 


The interactions between the magnetic fields, the jets, and the accretion disc 
are to slow down the protostar’s rotation, which would explain why most stars 
have a much slower spin than protostars of similar mass. 

Since the 1990s, the discovery of accretion discs around new stars led 
astronomers to speculate that these are the precursors to possible planetary 
formation. Many of these splendid objects were discovered in Orion but are, 
naturally, unobservable for the amateur astronomer. 


3.4 Clusters and Groups of Stars 


Stars do not form in isolation. 12 You do not get one star forming here, and 
perhaps another forming overthere! A dark nebulae can contain the material that 
could form hundreds of stars, and so stars tend to form in groups, or clusters. 

3.4.1 Galactic Star Clusters 

The stars that form out of the same cloud of material will not necessarily all 
have the same mass. Far from it — the masses will differ, and, as a consequence, 
reach the main sequence at different times. As I mentioned earlier, high-mass 
stars evolve faster than low-mass stars, and so at a time when these high-mass 
stars are shining brightly as stars in their own right, the low-mass protostars may 
still be cocooned within their dusty mantles. Consequently, the intense radiation 
emitted by the new, hot, and bright stars may disturb the normal evolution of 
the low-mass stars, and so reduce their final mass. 

Over time, however, the stellar nursery of young stars will gradually disperse. 
Calculations predict that massive stars have much shorter life-spans than smaller, 
less massive ones, so you can easily see that some stars (the more massive ones) 
do not live long enough to escape their birthplace, whereas a smaller star, say, 
of solar mass size, will in most cases easily escape from its stellar birthplace. 

It’s worth noting that, in relation to stars of mass about equal to that of the 
Sun, where there may be several thousand objects, the combined gravitational 
attraction of so many stars may slow down the dispersion of the group. It really 
depends on the star density and mass of the particular cluster. Thus, the most 
dense or closely packed clusters that contain solar-mass-sized stars will be the 
ones that contain the oldest population of stars, while the most open clusters 
will have the youngest star population. 

Open clusters, or galactic clusters, as they are sometimes called, are collections 
of young stars containing anywhere from a dozen members to hundreds. A few 
of them (for example, Messier 11 in Scutum) contain an impressive number of 
stars, equalling that of globular clusters, while others seem little more than a faint 
grouping set against the background star field. Such is the variety of open clusters 
that they come in all shapes and sizes. Several are over a degree in size and their 
full impact can only be seen by using binoculars, as a telescope has too narrow 
a field of view. An example of such a large cluster is Messier 44 in Cancer. Then 
there are tiny clusters, seemingly nothing more than compact multiple stars, as 
is the case with IC 4996 in Cygnus. In some cases, all the members of the cluster 




Stars 


73 


are equally bright, such as Caldwell 71 in Puppis; but there are others that consist 
of only a few bright members accompanied by several fainter companions, as is 
the case of Messier 29 in Cygnus. The stars that make up an open cluster are 
called Population I stars, which are metal-rich and usually found in or near the 
spiral arms of the Galaxy. 

The size of a cluster can vary from a few dozen l.y., as in the case of NGC 255 
in Cassiopeia, to about 70 l.y., as in either component of Caldwell 14, the Perseus 
Double Cluster. 

The reason for the varied and disparate appearances of open clusters is the 
circumstances of their births. It is the interstellar cloud that determines both 
the number and type of stars that are born within it. Factors such as the size, 
density, turbulence, temperature, and magnetic field all play a role as the deciding 
parameters in star birth. In the case of giant molecular clouds, or GMCs, the 
conditions can give rise to both O- and B-type giant stars along with solar-type 
dwarf stars — whereas in small molecular clouds (SMCs), only solar-type stars will 
be formed, with none of the luminous B-type stars. An example of an SMC is the 
Taurus Dark Cloud, which lies just beyond the Pleiades. 

By observing a star cluster, we can study in detail the process of star formation 
and interaction between low- and high-mass stars. As an example, look at 
Figure 3.5, which shows the H-R diagram for the cluster NGC 2264, located in 
Monoceros. Note that all the high-mass stars, which are the hottest stars, with a 
temperature of about 20,000 K, have already reached the main sequence, while 
those with temperatures at about 10,000 K or below have not. These low-mass 
and cooler stars are in the latter stages of pre-main-sequence star formation, with 
nuclear fusion just about beginning at their cores. Astronomers can compare 
this H-R diagram with the theoretical models, and they have deduced that this 
particular cluster is very young, only two million years old. 

This young star cluster is about 800 pc from Earth and contains many 
T Tauri stars. Each dot is a star whose temperature and luminosity have been 
measured. 

By comparison, we can look at the H-R diagram for a very famous cluster — the 
Pleiades star cluster, Figure 3.6. We can see right away that the cluster must 
be older than NGC 2264 because most of the stars are already on the main 
sequence. From studying the Pleiades H-R diagram, astronomers believe that 
the cluster could be about 50 million years old. Also look at the area on the 
H-R diagram with a temperature of about 10,500 K and luminosities ranging 
from 10 to 1O 2 L 0 . You will see a few stars that do not seem to he on the main 
sequence. This is not the case because they are still in the process of being 
formed; on the contrary, these massive stars have left the main sequence. They 
were among the first to be formed and thus are the oldest, and they are now 
evolving into a different kind of star. As we shall see later, all the hydrogen 
at the center of these stars has been used up, 13 and helium-burning is now 
proceeding. 

This is a much older cluster, around 50 million years. Most, if not all, of 
the low-mass cool stars have reached the main sequence, which implies that 
hydrogen-burning has started in their cores. 

An interesting aspect of open clusters is their distribution in the night sky. 
You may be forgiven in thinking that they are randomly distributed across the 


74 


Astrophysics is Easy 


Luminosity 

(L.) 



Figure 3.5. H-R diagram of star cluster NGC 2264. 


sky, but surveys show that although well over a thousand clusters have been 
discovered, only a few are observed to be at distances greater than 25° above or 
below the galactic equator. Some parts of the sky are very rich in clusters — e.g., 
Cassiopeia and Puppis — and this is due to the absence of dust lying along these 
lines of sight, allowing us to see across the spiral plane of our Galaxy. Many of 
the clusters mentioned here actually lie in different spiral arms, and so as you 
observe them, you are actually looking at different parts of the spiral structure 
of our Galaxy. 

I mentioned earlier that stars are not born in isolation. Nor are they born 
simultaneously. Recall that the more massive a star, the faster it contracts and 
becomes stable, thus joining the main sequence; this results in some clusters’ 
having bright young O and B main-sequence stars, while at the same time 
containing low-mass members, which may still be in the process of gravitational 
contraction (for example, the star cluster at the center of the Lagoon Nebula). In 
a few cases, the star production in a cluster is at a very early stage, with only a 



Stars 


75 



Figure 3.6. H-R diagram of the Pleiades star cluster. 


few stars visible; the majority are still in the process of contraction and hidden 
within the interstellar cloud. 

A perfect example of such a process is the open cluster within Messier 42, 
the Orion Nebula. The stars within the cluster, the Trapezium, are the brightest, 
youngest, and most massive stars in what will eventually become a large cluster 
containing many A-, F-, and G-type stars. However, the majority are blanketed 
by the dust and gas clouds and are only detectable by their infrared radiation. 

As time passes, the dust and gas surrounding a new cluster will be blown away 
by the radiation from the O-type stars, resulting in the cluster’s becoming visible 
in its entirety, such as in the case of the young cluster Caldwell 76 in Scorpius. 

Once a cluster has formed, it will remain more or less unchanged for at least a 
few million years, but then changes within the cluster may occur. Two processes 
are responsible for changes within any given cluster. The evolution of open 
clusters depends on both the initial stellar content of the group and the ever- 
pervasive pull of gravity. If a cluster contains 0-, B-, and A-type stars, these stars 



76 


Astrophysics is Easy 


will eventually become supernovae, leaving the cluster with slower-evolving, less- 
massive, and less-luminous members of type A and M stars. A famous example 
of such a cluster is Caldwell 94, the Jewel Box in Crux, which is the highlight 
of the southern sky, and, alas, unobservable to northern hemisphere observers. 
However, these, too, will become supernovae, with the result that the most 
luminous members of a cluster will, one by one, disappear over time. This does 
not necessarily mean the demise of a cluster, especially those that have many tens 
or hundreds of members. But some, which consist of only a few bright stars, will 
seem to meld into the background star field. However, even those clusters that 
have survived the demise of their brighter members will eventually begin to feel 
the effect of a force that pervades everywhere — the Galaxy’s gravitational field. As 
time passes, the cluster will be affected by the influence of other globular clusters 
and the interstellar matter itself, as well as the tidal force of the Galaxy. The 
cumulative effect of all of these encounters will be that some of the less-massive 
members of the cluster acquire enough velocity to escape from the cluster. Thus, 
given enough time, a cluster will fade and disperse. (Take heart, as this is not 
likely to happen in the near future so that you would notice: the Hyades star 
cluster, even after having lost most of its K- and M-type dwarf stars, is still with 
us after 600 million years!). 

For the amateur, observing open clusters is a very rewarding experience, as 
they are readily observable, from naked-eye clusters to those visible only in 
larger telescopes. Many of them are best viewed with binoculars, especially the 
larger clusters that are of an appreciable angular size. Furthermore, nearly all 
have double or triple stars within the cluster, and so, regardless of magnification, 
there is always something interesting to be seen. 

From the preceding chapter, you know that color in observed stars is best 
seen when contrasted with a companion(s). Thus, an open cluster presents a 
perfect opportunity for observing star colors. Many clusters, such as Pleiades, 
are all a lovely steely blue color. On the other hand, Caldwell 10 in Cassiopeia 
has contrasting bluish stars along with a nice orange star. Other clusters have 
a solitary yellowish or ruddy orange star along with fainter white ones, such 
as Messier 6 in Scorpius. An often-striking characteristic of open clusters is 
the apparent chains of stars that are seen. Many clusters have stars that arc 
across apparently empty voids, as in Messier41 in Canis Major. 

Because open clusters display such a wealth of characteristics, different param- 
eters are assigned to a cluster that describe its shape and content. For instance, a 
designation called the Trumpler type is often used. It is a three-part designation 
that describes the cluster’s degree of concentration — that is, from a packed cluster 
to one that is evenly distributed, the range in brightness of the stars within the 
cluster, and finally the richness of the cluster from poor (fewer than 50 stars) to 
rich (more than 100). The full classification is: 

Trumpler Classification for Star Clusters 

Concentration 

I Detached — strong concentration of stars toward the center. 

II Detached — weak concentration of stars toward the center. 

III Detached — no concentration of stars toward the center. 

IV Poor detachment from background star field. 


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77 


Range of brightness 

1 Small range. 

2 Moderate range. 

3 Large range. 

Richness of cluster 

p Poor (with fewer than 50 stars), 
m Moderate (with 50-100 stars), 
r Rich (with more than 100 stars), 
n Cluster within nebulosity. 

Two points that can often cause problems need to be mentioned: the magnitude 
and size of the cluster. The quoted magnitude of a cluster may be the result of 
only a few bright stars, or, on the other hand, may be the result of a large number 
of faint stars. Also, the diameter of a cluster is often misleading, as in most cases 
it has been calculated from photographic plates, which, as experienced amateurs 
will know, bear little resemblance to what is seen at the eyepiece. 

Although magnitudes and diameters may be quoted in the text, do treat them 
with a certain amount of caution. 

In the descriptions given below, the first line lists the name, the position, 
and the approximate midnight transit time; the second line presents the visual 
magnitude (the combined magnitude of all stars in cluster), the object size in 
arcminutes ©, the approximate number of stars in the cluster (bear in mind that 
the number of stars seen will depend on magnification and aperture and will 
increase when large apertures are used, thus the number quoted is an estimate 
using modest aperture), the Trumpler designation, and the level of difficulty 
(based on the magnitude, size, and ease of finding the cluster). 


3.4.1 .1 Bright Star Clusters 

Messier 41 NGC 2287 06 h 47.0 m -20° 44' Dec-Jan-Feb 

4.5m © 38^ 70 113 m Canis Major 

Easily visible to the naked eye on clear nights as a cloudy spot slightly larger 
in size than the full moon. Contains blue B-type giant stars, as well as several 
K-type giants. Current research indicates that the cluster is about 100 million 
years old and occupies a volume of space 80 l.y. in diameter. 

Caldwell 64 NGC2362 07 h 18.8 m -24°57' Dec-Jan-Feb 

4.1m ©8 r 60 13 pn Canis Major 

A very nice cluster, tightly packed, and easily seen with small binoculars. The glare 
from rCMa tends to overwhelm the majority of stars, but the cluster becomes 
truly impressive with telescopic apertures; the bigger the aperture, the more 
stunning the vista. It is believed to be very young — only a couple of million years 


78 Astrophysics is Easy 

old — and thus has the distinction of being the youngest cluster in our Galaxy. 
Contains O- and B-type giant stars. 

Messier 48 NGC 2548 08 h 13.8 m -05°48' Dec-Jan-Feb 

5.8m ©55' 80 13 r Hydra 

Located in a rather empty part of the constellation Hydra, this is believed to be the 
missing Messier object. It is a nice cluster and can be viewed both in binoculars 
and small telescopes. In the former, about a dozen stars are seen, with a pleasing 
triangular asterism at its center, while the latter will show a rather nice but large 
group of about 50 stars. Many amateurs often find the cluster difficult to locate for 
the reason mentioned above, but also for the fact that within a few degrees of M48 is 
another nameless, but brighter, cluster of stars, which is often mistakenly identified 
as M48. Some observers claim that this nameless group of stars is in fact the correct 
missing Messier object, and not the one which now bears the name. 

Messier 44 NGC 2632 08 h 40.1 m +19°59' Dec-Jan-Feb 

3.1m ©95' 60 II2m Cancer 

A famous cluster called Praesepe (the Manger) or the Beehive. One of the largest 
and brightest open clusters from the viewpoint of an observer. An old cluster, 
about 700 million years, distance 500 l.y., with the same space motion and 
velocity as the Hyades, which suggests a common origin for the two clusters. A 
nice triple star, Burnham 584, is located within M44, located just south of the 
cluster’s center. A unique Messier object in that it is brighter than the stars of 
the constellation within which it resides. Due to its large angular size in the sky, 
it is best seen through binoculars or a low-power eyepiece. 

Caldwell 54 NGC 2506 08 h 00.2 m -10°47' Dec-Jan-Feb 

7.6m ©7' 100 12 r Monoceros 

A nice rich and concentrated cluster best seen with a telescope, but one that is 
often overlooked due to its faintness, even though it is just visible in binoculars. 
Includes many 11th- and 12th-magnitude stars. It is a very old cluster, about 
2 billion years, and contains several blue stragglers. These are old stars that 
nevertheless have the spectrum signatures of young stars. This paradox was 
solved when research indicated that the young-looking stars are the result of a 
merger of two old stars. 

Messier 67 NGC 2682 08 h 50.4 m +11°49' Jan-Feb-Mar 

6.9m ©30' 200 II2m Cancer 

Often overlooked because of its proximity to M44, it is nevertheless very pleasing. 
However, the stars which it is composed of are faint, and so in binoculars it 
will be unresolved and seen as a faint misty glow. At a distance of 2500 l.y. it is 
believed to be very old, possibly 9 billion years, and thus has had time to move 
from the Galactic Plane, the usual abode of open clusters, to a distance of about 
1600 l.y. off the plane. 

Caldwell 76 NGC 6231 16 h 54.0 m -41°48' May-Jun-Jul 
2.6m © 14' 100 13 p Scorpius 


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79 


A superb cluster located in an awe-inspiring region of the sky. Brighter by 2.5 
magnitudes than its northern cousins, the double cluster in Perseus. The cluster 
is full of spectacular stars: very hot and luminous O-type and BO-type giants and 
supergiants, a couple of Wolf-Rayet stars, and £ _1 Scorpii, which is a B1.5 la 
extreme supergiant star with a luminosity nearly 280,000 times that of the Sun! 
The cluster is thought to be a member of the stellar association 14 Sco OBI, with 
an estimated age of 3 million years. A wonderful object to view in binoculars 
and telescopes, the cluster contains many blue, orange, and yellow stars. It lies 
between p} +1 Scorpii and £ _1 Scorpii, an area rich in spectacular views. 

Trumpler 24 Harvard 12 16 h 57.0 m -40° 40' May-Jun-Jul 
8.6m © 60' 100 IV2 pn Scorpius 

A loose and scattered cluster, set against the backdrop of the Milky Way. It is, 
along with nearby Collinder 316, the core of the Scorpius OBI stellar association. 

Messier 7 NGC 6475 17 h 53.9 m -34° 49' May-Jun-Jul 
3.3m ©80' 80 13 r Scorpius 

An enormous and spectacular cluster. It presents a fine spectacle in binoculars 
and telescopes, containing more than 80 blue-white and pale yellow stars. It is 
only just over 800 l.y. away, but it is over 200 million years old. Many of the 
stars are around 6th and 7th magnitude and thus should be resolvable with the 
naked eye. 

Messier 24 15 18 h 16.5 m -18°50' May-Jun-Jul 

2.5m ©95'x35' Sagittarius 

Another superb object for binoculars. This is the Small Sagittarius Star Cloud, 
visible to the naked eye on clear nights, and nearly four times the angular size 
of the Moon. The cluster is in fact part of the Norma Spiral Arm of our Galaxy, 
located about 15,000 l.y. from us. The faint background glow from innumerable 
unresolved stars is a backdrop to a breathtaking display of 6th- to lOth-magnitude 
stars. It also includes several dark nebulae which adds to the three-dimensional 
impression. Many regard the cluster as truly a showpiece of the sky. 

Messier 16 NGC 6611 18 h 18.8 m -13°47' May-Jun-Jul 

6.0m ©22' 50 113 mn Serpens Cauda 

A fine large cluster easily seen with binoculars. It is about 7000 l.y. away, located 
in the Sagittarius-Carina Spiral Arm of the Galaxy. Its hot O-type stars provide 
the energy for the Eagle Nebula, within which the cluster is embedded. 

Messier 25 IC 4725 18 h 31.6 m -19°15' May-Jun-Jul 
4.6m © 32' 40 13 m Sagittarius 

Visible to the naked eye, this is a pleasing cluster suitable for binocular obser- 
vation. It contains several star chains and is also noteworthy for small areas of 
dark nebulosity that seem to blanket out areas within the cluster, but you will 


80 


Astrophysics is Easy 


need perfect conditions to appreciate these. Unique for two reasons: it is the only 
Messier object referenced in the Index Catalogue (IC), and it is one of the few 
clusters to contain a cepheid- type variable star — U Sagittarii. The star displays a 
magnitude change from 6.3 to 7.1 over a period of 6 days and 18 hours. 

Messier 11 NGC 6705 18 h 51.1 m -06°16' Jun-Jul-Aug 
5.8m © 13' 200 12 r Scutum 


Also known as the Wild Duck Cluster, this is a gem of an object. Although 
it is visible with binoculars as a small, tightly compact group reminiscent of 
a globular cluster, they do not do it justice. With telescopes, its full majesty 
becomes apparent. Containing many hundreds of stars, it is a very impressive 
cluster. It takes high magnification, as well, where many more of its 700 members 
become visible. At the top of the cluster is a glorious pale yellow-tinted star. 

IC 1396 21 h 39.1 m +57°30' Jul-Aug-Sep 

3.7m © 50' 40 II mn Cepheus 

Although a telescope of at least 20 cm is needed to truly view this cluster, it is 
nevertheless worth searching for. It lies south of Herschel’s Garnet Star and is 
rich but compressed. What makes this so special, however, is that it is cocooned 
within a very large and bright nebula. 

Caldwell 13 NGC 457 01 h 19.1 m +58°20 Sep-Oct-Nov 
6.4m © 13' 80 13 r Cassiopeia 


This is a wonderful cluster and can be considered as one of the finest in Cassiopeia. 
Easily seen in binoculars as two southward-arcing chains of stars, surrounded by 
many fainter components. The gorgeous blue and yellow double, cp Cass, and a 
lovely red star, HD 7902, lie within the cluster. Located at a distance of about 
8000 l.y., this young cluster is located within the Perseus Spiral Arm of our Galaxy. 


Caldwell 14 NGC 869 02 h 19.0 m 

5.3m ©29' 200 

NGC 884 02 h 22.4 m 

6.1m ©29' 115 


+57°09' 
13 r 

+57°07' 

112 p 


Sep— Oct— Nov 
Perseus 


The famous Double Cluster in Perseus is a highlight of the northern-hemisphere 
winter sky. Strangely, never catalogued by Messier. Visible to the naked eye 
and best seen using a low-power, wide-field optical system. But whatever system 
is used, the views are marvellous. NGC 869 has about 200 members, while 
NGC 884 has about 150. Both are composed of A-type and B-type supergiant 
stars with many nice red giant stars. However, the systems are dissimilar; NGC 
869 is 5.6 million years old (at a distance of 7200 l.y.), whereas NGC 884 is 
younger at 3.2 million (at a distance of 7500 l.y.). But be advised that in astro- 
physics, distance and age determination are error prone! Also, it was found 
that nearly half the stars are variables of the type B, indicating that they are 
young stars with possible circumstellar discs of dust. Both are part of the Perseus 
OBI Association 16 from which the Perseus Spiral Arm of the Galaxy has been 


Stars 


81 


named. Do not rush these clusters, but spend a long time observing both of them 
and the background star fields. 

Messier 45 Melotte 22 03 h 47.0 m +24°07' Oct— Nov— Dec 

1.2m ©110' 100 13 r Taurus 

Without a doubt, the sky’s premier star cluster. The Seven Sisters or Pleiades 
is beautiful however you observe it — naked eye, through binoculars, or with a 
telescope. To see all the members at one go will require binoculars or a rich-field 
telescope. Consisting of more than 100 stars, spanning an area four times that 
of the full Moon, it will never cease to amaze people. It is often stated that from 
an urban location, 6 to 7 stars may be glimpsed with the naked eye. However, 
it may come as a surprise to many of you that it has 10 stars brighter than 6th 
magnitude, and that seasoned amateurs with perfect conditions have reported 
18 being visible with the naked eye. It lies at a distance of 410 l.y., is about 
20 million years old (although some report it as 70 million), and is the 4th- 
nearest cluster. Messier 45 contains many stunning blue and white B-type giants. 
The cluster contains many double and multiple stars. Under perfect conditions 
with exceptionally clean optics, the faint nebula NGC 1435, the Merope Nebula 
surrounding the star of the same name ( Merope - 23 Tauri), can be glimpsed and 
was described by W. Tempel in 1859 as “a breath on a mirror.” However, this 
and the nebulosity associated with the other Pleiades are not, as they were once 
thought to be the remnants of the original progenitor dust and gas cloud. The 
cluster is just passing through an edge of the Taurus Dark Cloud Complex. It is 
moving through space at a velocity of about 40 kilometers a second, so by 32,000 
AD, it will have moved an angular distance equal to that of the full Moon. The 
cluster contains the stars Pleione, Atlas, Alcyone, Merope, Maia, Electra, Celaeno, 
Taygeta, and Asterope. A true celestial showpiece. 

Caldwell 41 Melotte 25 04 h 27.0 m +16°00' Oct-Nov-Dec 

0.5m ©330' 40 113 m Taurus 

Also known as the Hyades. The nearest cluster after the Ursa Major Moving 
Stream, lying at a distance of 151 l.y. with an age of about 625 million years. 
Even though the cluster is widely dispersed both in space and over the sky, it 
nevertheless is gravitationally bound, with the more massive stars lying at its 
center. Best seen with binoculars due to the large extent of the cluster — over 
5 i,,2 °. Hundreds of stars are visible, including the fine orange giant stars y, S, 
e, and 0~ l Tauri. Aldebaran, the lovely orange K-type giant star, is not a true 
member of the cluster, but it is a foreground star only 70 l.y. away. Visible even 
from light-polluted urban areas — a rarity! 

Collinder 69 - 05 h 35.1 m +09°56' Nov— Dec— Jan 

2.8m © 65' 20 113 pn Orion 

This cluster surrounds the 3rd-magnitude star A Orionis and includes <p~ l and 
( p~ 2 Orionis, both 4th-magnitude. Encircling the cluster is the very faint emission 


82 Astrophysics is Easy 

nebula Sharpless 2-264, only visible using averted vision and an OIII filter with 
extremely dark skies. Perfect for binoculars. 

Messier 37 NGC 2099 05 h 52.4 m +32°33' Nov-Dec-Jan 
5.6m ©20' 150 III r Auriga 

The finest cluster in Auriga. Contains many A-type stars and several red giants. 
Visible at all apertures, from a soft glow with a few stars in binoculars to a fine, 
star-studded field in medium-aperture telescopes. In small telescopes using a low 
magnification, it can appear as a globular cluster. The central star is colored a 
lovely deep red, although several observers report it as a much paler red, which 
may indicate that it is a variable star. Visible to the naked eye. 

Collinder 81 NGC 2158 06 h 07.5 m +24°06' Nov-Dec-Jan 
8.6m ©5' 70 113 r Gemini 

Lying at a distance of 16,000 l.y., this is one of the most distant clusters visible 
using small telescopes, and lying at the edge of the Galaxy. It needs a 20 cm 
telescope to be resolved, and even then only a few stars will be visible against a 
background glow. It is a very tight, compact grouping of stars, and something 
of an astronomical problem. Some astronomers class it as intermediate between 
an open cluster and a globular cluster, and it is believed to be about 800 million 
years old, making it very old as open clusters go. 


3.4.2 Stellar Associations and Streams 

There exists another type of grouping of stars, which is much more ephemeral 
and spread over a large region of the sky, and although not strictly associated 
with star formation, it is, however, an integral part of star evolution. Furthermore, 
as this book is dealing with both the evolution and observational properties of 
stars, I think it wise to mention it here. 

A stellar association is a loosely bound group of very young stars. They may 
still be swathed in the dust and gas cloud within which they formed, and star 
formation may still be occurring within the cloud. They differ from open clusters 
in their enormity, covering both a sizable angular area of the night sky and at the 
same time encompassing a comparably large volume in space. As an illustration 
of this huge size, the Scorpius-Centaurus Association is about 700 by 760 l.y. in 
extent, and it covers about 80°. 

There are three types of stellar associations: 

OB associations, containing very luminous O- and B-type main-sequence, giant, 
and supergiant stars. 

B associations, containing only B-type main-sequence and giant stars but with 
an absence of O-type stars. These associations are just older versions of the 
OB association, and thus the faster-evolving O-type stars have been lost to the 
group as supernovae. 


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T associations, which are groupings of T Tauri-type stars. These are irregular 
variable stars that are still contracting and evolving toward being A-, F-, and 
G-type main-sequence stars. As they are still in their infancy, more often than 
not they will be shrouded in dark dust clouds, and those that are visible will 
be embedded in small reflection and emission nebulae (see Chapter 4). 

The OB associations are truly enormous objects, often covering many hundreds 
of l.y. This is a consequence of the fact that massive O- and B-type stars can 
only be formed within the giant molecular clouds which are themselves hundreds 
of l.y. across. On the other hand, the T associations are much smaller affairs, 
perhaps only a few l.y. in diameter. In some cases, the T association is itself 
located within or near an OB association. 

The lifetime of an association is comparatively short. The very luminous O-type 
stars are soon lost to the group as supernovae, and, as usual, the ever-pervasive 
gravitational effects of the Galaxy soon disrupt the association. The coherence and 
identity of the group can only exist for as long as the brighter components stay 
in the same general area of a spiral arm, as well as having a similar space motion 
through the Galaxy. As time passes, the B-type stars will disappear through stellar 
evolution, and the remaining A-type and later stars will now be spread over 
an enormous volume of space, and the only common factor among them will 
be their motion through space. At that point, the association is called a stellar 
stream. An example of such a stream and one that often surprises the amateur 
(it did me!) is the Ursa Major Stream. This is an enormous group of stars, with 
five central stars of Ursa Major (The Plough) being its most concentrated and 
brightest members. The stream is also known as The Sirius Supercluster after its 
brightest member. The Sun actually lies within this stream (more information 
about this fascinating stream can be found below). 

3.4.2. 1 Bright Stellar Associations and Streams 

The Orion Association, 1600 l.y. 

This association includes most of the stars in the constellation down to 3.5 
magnitude, except for y Orionis and 77 3 Orionis. Also included are several 4th-, 
5th-, and 6th-magnitude stars. The wonderful nebula M42 is also part of this 
spectacular association. Several other nebulae (including dark, reflection, and 
emission nebulae) are all located within a vast Giant Molecular Cloud, which is 
the birthplace of all the O- and B-type supergiant, giant, and main-sequence stars 
in Orion. The association is believed to be 800 l.y. across and 1000 l.y. deep. By 
looking at this association, you are in fact looking deep into our own spiral arm, 
which, incidentally, is called the Cygnus-Carina Arm. 


The Scorpius-Centaurus Association, 550 l.y. 

A much older, but closer association than the Orion association. It includes most 
of the stars of 1st, 2nd, and 3rd magnitude in Scorpius down through Lupus and 
Centaurus to Crux. Classed as a B-type association because it lacks O-type stars, 


84 


Astrophysics is Easy 


its angular size on the sky is about 80°. It is estimated to be 750 x 300 l.y. in size, 
and 400 l.y. deep, with the center of the association midway between a Lupi and 
£ Centauri. Its elongated shape is thought to be the result of rotational stresses 
induced by its rotation about the Galactic center. Bright stars in this association 
include Ophiuchi, j3, v, 8 and cr Scorpii, a, y Lupi, e, 8, /jl and e Centauri, and fi 
Crucis. 


The Zeta Persei Association, 1300 l.y. 

Also known as Per OB2, this association includes £ and £ Persei, as well as 40, 
42, and o Persei. The California nebula, NGC 1499, is also within this association. 


The Ursa Major Stream, 75 l.y. 

As briefly mentioned earlier in this section, this stream includes the five central 
stars of the Plough. It is spread over a vast area of sky, approximately 24°, and 
is about 20 x 30 l.y. in extent. It includes as members Sirius ( a Canis Majoris), 
a Coronae Borealis, 8 Leonis, (3 Eridani, 8 Aquarii, and (3 Serpentis. Due to 
the predominance of A1 and A0 stars within the association, its age has been 
estimated at 300 million years. 


The Hyades Stream 

There is some evidence (although it is not fully agreed upon) that the Ursa Major 
stream is itself within a much older and larger stream. This older component 
includes M44, Praesepe in Cancer, and the Hyades in Taurus, with these two 
open clusters being the core of a very large but loose grouping of stars. Included 
within this are Capella (a Aurigae), a Canum Venaticorum 1 , 8 Cassiopeiae, and 
A Ursae Majoris. The stream extends to over 200 l.y. beyond the Hyades star 
cluster, and 300 l.y. behind the Sun. Thus, the Sun lies within this stream. 1 


The Alpha Persei Stream 540 l.y. 

Also known as Melotte 20, this is a group of about 100 stars, including a Persei, 
if/ Persei, 29, and 34 Persei. The stars 8 and e Persei are believed to be among its 
most outlying members, as they also share the same space motion as the main 
groups of stars. The inner region of the stream is measured to be over 33 l.y. in 
length, the distance between 29 and if/ Persei. 


3.5 Star Formation Triggers 


We have seen how stars are formed from clouds of dust and gas, and how these 
clouds clump together under the force of gravity to form protostars. In addition, 
the evolution of a protostar to a main sequence depends on the initial mass of 



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the protostar and so determines where it will arrive on the main sequence. But 
one thing we have not yet addressed is what causes a protostar to form in the 
first place! This is the topic of the final part of this section. 

The mechanisms that provide the “triggers” for star formation have three very 
disparate origins: 

• The spiral arms of a galaxy 

• Expanding HII regions 

• Supernovae 

We mentioned earlier in this section that the spiral arms of galaxies are a prime 
location for star formation because the gas and dust clouds temporarily “pile up” 
as they orbit around the center of a galaxy. 17 In such a spiral arm, the molecular 
clouds are compressed as it passes through the region. In the molecular cloud’s 
densest regions, vigorous star formation can then occur. 

Massive stars, such as O-type and B-type stars, emit immense amounts of 
radiation, usually in the ultraviolet part of the spectrum. This, in turn, causes 
the surrounding gas to ionize, and an HII region is formed within the larger 
molecular cloud. The strong stellar winds and ultraviolet radiation that O- and 
B-type stars possess can carve out a cavity within the molecular cloud into which 
the HII region expands. The stellar wind is moving at such a high velocity that 
it is supersonic (i.e., faster than the speed of sound in that particular region). A 
shock wave associated with this supersonically expanding HII region then collides 
with the rest of the molecular cloud. In doing so, it compresses the cloud, and so 
further star formation occurs. The new O- and B-type stars that result from this 
induce further star formation, but at the same time, the precursor O- and B-type 
stars that originally started the procedure may well have dispersed by this point. 
In this manner, an OB association “devours” a molecular cloud, leaving older 
stars in its wake. 

The Orion Nebula is one example of such a mechanism, where the four stars 
of the Trapezium are ionizing the surrounding material. The nebula itself is at 
the edge of a giant molecular cloud, some 500, 000 M Q . 

The final mechanism which is believed to induce further star formation is a 
supernova. As we shall see in later sections, a supernova is the death of a star 
and results in a catastrophic explosion, usually blowing the star to bits! What 
is important to us at this stage is that the outer layers of the star are ejected 
into space at incredible speeds, may be several thousand kilometers per second! 
This shock wave, an expanding shell of material, will be moving at supersonic 
velocities and, in a similar manner as mentioned above, will impact the material 
in the interstellar medium, and in doing so will compress and heat it. In doing 
so, it will stimulate further star formation. 

We have now covered the amazing processes involved in star formation, from 
vast clouds of dust and gas to glowing spheres of nuclear fusion — the birth of 
a star. However, do not think that we know all there is to know about star 
birth because we do not! For instance, a spiral arm that passes through a giant 
molecular cloud tends to produce giant O- and B-type stars, whereas the stars 
induced by supernovae shock waves are predominantly A-, F-, G-, and K-type 
stars. Also, in our home Galaxy, there often seems to be a lot of dust associated 


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with star formation that shields the newly born stars from the destructive effects 
of ultraviolet radiation from other hot stars that are close by. However, in a 
nearby galaxy — The Large Magellanic Cloud 18 (LMC) — it has been observed that 
young OB associations have hardly any dust at all! Nevertheless, what we do 
know is amazing and involves mechanisms from star death to the rotation of 
galaxies. 

Most of the stars that we observe in the night sky have one thing in common: 
they are on the main sequence. There are, of course, exceptions. Betelgeuse has 
left the main sequence and has become a red giant star; the hydrogen-burning 
at the center of its core has stopped, and now helium is burning by fusion 
processes. Also, Sirius B has evolved far from the main sequence and has become 
a white dwarf star, with no nuclear fusion occurring within it. But for the large 
majority, the main sequence is a stable time, with only small changes in mass and 
luminosity occurring. However, as a star ages, changes occur in the way energy 
is formed, and this in turn affects its size and thus its luminosity, and so the 
star leaves the main sequence to begin the next phase of its life. The remaining 
material in this chapter will look at these periods in a star’s life, whether it is a 
small, low-mass, and cool star or a bright, high-mass, and hot star. 

Before we look at the many types of stars on the main sequence, it will be 
helpful (and indeed necessary) for us to look at the nearest star to us — the Sun. 
After all, astronomers have been studying the closest star to us for a long time 
now, and so we have a good idea of what is going on . 19 In looking at the Sun in 
detail, we will be able to see how energy is produced in the core and how this 
energy is transported to the surface, and then to us on the Earth! We can then 
look at other stars and compare and contrast them with what we know about 
the Sun. 


3.6 The Sun— The Nearest Star 


In this section, we shall look at the Sun bearing in mind it is a star on the main 
sequence. Therefore, I shall not discuss in any depth topics such as sunspots, the 
sunspot cycle, and so on . 20 Instead, we will concentrate on the internal structure, 
the means of energy production, and the manner in which energy is transported 
from its source to us on Earth. With this approach, it is possible to use the Sun 
as a benchmark with which to compare stars of differing size. 

Due to the advances not only in astronomy but in computing as well, 
astronomers have been able to determine the conditions inside the Sun by 
solving several equations that describe how the temperature, mass, luminosity, 
and pressure change with distance from the center of the Sun. To solve them, 
we need to know the mechanisms by which energy is transported throughout 
the Sun, either by radiation or convection, the chemical composition of the Sun, 
and the rate of energy production at any specific distance from its center. Now 
although the equations are simple to solve, computers are needed, so we will just 
say that the results seem to match the observations, which is always a good test 
for any theory. 




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3.6.1 From the Core to the Surface 


The Sun’s internal structure is shown in Figure 3.7. The visible surface of the 
Sun, called the photosphere, has a temperature of about 5800 K, and although it 
may look like a well-defined surface from the Earth, it is in fact a gas less dense 
than the Earth’s atmosphere. Both the density and temperature increase steadily 
as we progress from the surface to the core. Beneath the photosphere is a very 
turbulent area called the convection zone, where energy generated in the core 
travels upward, transported by rising columns of hot gas and the falling of cool 
gas. This process is called convection. So the photosphere is in fact the top of the 
convection zone. Descending deeper through the convection zone, the pressure 
and density increase quite substantially, along with the temperature. The density 
is far greater than that of water, but remember that we are still talking about 
gas, albeit one in a very strange state. A gas under these extreme conditions of 
temperature and/or pressure is usually called plasma. 21 The temperature in this 
region is about 2 million K, and the solar plasma absorbs the photons. 

About one-third of the way down to the center, the very turbulent convection 
zone gives way to the more stable plasma of the radiation zone. Here the energy 
is transported outward primarily by photons of X-ray radiation. The temperature 
in this region is now about 10,000,000 K. At the central region, the core of the 
Sun, the temperature is now 1 5,000,000 K, and it is here that hydrogen is being 
transformed to helium. The pressure in this region is nearly 200 billion times the 
surface pressure found on the Earth. The central temperature and pressure are 
both impressive, with the core compressed to a density of about 150, 000 kg m 3 , 
which is about 150 times the density of water. It may come as a surprise to 
some people that essentially all of the Sun’s energy is produced in the inner 
25% of its radius, which corresponds to about 1.5% of its volume. This is a 
consequence of the very acute temperature sensitivity of nuclear reactions. If we 
were to actually go to a point about 1/4 the distance from the center of the Sun 
to its surface, the temperature would have fallen to about 8,000,000 K, and at 


1009 

70 



Nuclear reactions 
Core produce energy in core 


Energy produced in core 
is transported outward by 
photons 


Convection carries 
energy outward 



Nuclear-burning zone 


Figure 3.7. The internal structure of the Sun. 



88 


Astrophysics is Easy 


this lower temperature nuclear fusion energy production would have fallen to 
practically zero. So, virtually no energy is produced beyond the inner 25% of the 
solar radius. 

At the surface of the Sun, each kilogram of gas contains about 71% hydrogen, 
whereas in the core, the percentage of hydrogen will be much lower, around 34%. 
The reason for this is obvious — hydrogen has been the fuel for nuclear fusion 
for the past 4.6 billion years. The total power output of the Sun, which is its 
luminosity, is a staggering 3.8 x 10 26 joules per second. This may not mean much 
to most of us, but if we could somehow capture all of this energy, for even one 
second, it would be sufficient to meet all current energy demands for the human 
race for the next 50,000 years! But remember, only a tiny fraction of this reaches 
the Earth, as it is all dispersed in all directions into space. 

The current model of energy production in the Sun is that in which nuclear 
fusion is the generator. It is a source so efficient that the Sun will shine for 10 
billion years, and as it is only 4.6 billion years old at the moment, it has a long way 
to go! This current model of solar-energy generation means that the Sun’s size 
will generally be stable, maintained by a balance between the competing forces 
of gravity pulling inward and pressure pushing outward. This balance between 
forces is called hydrostatic equilibrium (or sometimes gravitational equilibrium). 
What this means is that at any given point within the Sun, the weight of the 
overlying material is supported by the underlying pressure. You may think that 
this is a simple concept, and so it is, but it maintains the integrity of the Sun, as 
well as most stars in the universe. When one or the other of the forces gains the 
upper hand, however, the consequences are spectacular, as we shall see in a later 
section. The hydrostatic equilibrium in the Sun means the pressure increases 
with the depth; this makes the Sun extremely hot and dense in its core. 

The efficiency with which the energy is transported outward by radiation is 
strongly influenced by the opacity of the gas through which the photons flow. 
The opacity describes the ability of a substance to stop the flow of photons. For 
instance, when the opacity is low (think of it as a clear day), photons are able 
to travel much greater distance between emission and re-absorption than when 
the opacity is high (a foggy, hazy day). If opacity is low, the transportation of 
energy by photons is very efficient. But when the opacity is high, the efficiency 
is reduced, which leads to an inefficient flow of energy and a higher rate of 
temperature decline. 


3.6.2 The Proton-Proton Chain 

To explain the Sun’s energy, we need a process that involves the most abundant 
element in the Sun, hydrogen. The fusion of hydrogen into helium was first 
proposed in 1920 by the British astronomer A. S. Eddington, although the details 
were not fully understood until 1940. 

Hydrogen, which is the lightest element, has a nucleus consisting of just one 
proton. The nucleus of helium, however, has four nuclear particles — two protons 
and two neutrons. So four hydrogen nuclei are needed to make one helium 
nucleus. But we cannot expect four protons to collide and instantly make a 
helium nucleus. This is unlikely to happen, as it has never happened before — not 


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even once — in the history of the universe. What happens instead is a series of 
reactions involving two reactions at a time. This series of reactions is called 
the proton-proton chain. 22 The reactions begin with an interaction between two 
protons that must come within 10 -15 meters of each other for a nuclear reaction 
to occur. There is a small problem, however, as protons are positively charged, 
and, just like magnets, repel each other. The result of this mutual repulsion is 
that most collisions between protons do not result in any reaction; instead, the 
two protons deflect each other and move apart. At room temperature, there is 
absolutely no possibility that two protons would collide with enough energy to 
get close enough to instigate a reaction. 

So, for any reactions to occur, we need conditions that will allow protons to 
move at very high velocities; these conditions exist in the center of the Sun (and, 
of course, stars). At the core of the Sun, the temperature is 15 million K, and a 
typical proton will be travelling at about 1 million kilometers per hour. But even 
at this fantastic speed, the likelihood of a reaction occurring is very small. If we 
could watch a single proton to see how long it would take before it eventually 
reacted with another proton in nuclear fusion, we would be waiting for about 5 
billion years! The important point here is that there are so many protons in the 
Sun’s core, every second 10 34 of them can undergo a reaction. 

The sequence of steps in the proton-proton chain is shown in Figure 3.8. 

Step 1: Two protons fuse to form a nucleus consisting of one proton and one 
neutron. This is the isotope of hydrogen called deuterium ( 2 H). The other 
products formed are a positively charged electron, called & positron ((3 ' ), 
and a neutrino (v), a minuscule particle with a tiny mass. The positron 


H 


H 


O 


O 



H 


Figure 3.8. The reactions of the proton-proton chain. 



Astrophysics is Easy 

does not last long, however; it soon meets up with an ordinary electron, 
and the result is the creation of two gamma rays that are rapidly absorbed 
by the surrounding gas, which consequently heats up. What happens to 
the neutrino? We shall discuss that later. 

The deuterium now fuses with a proton, producing a helium nucleus 
( 3 He) and gamma rays. The 3 He nucleus consists of two protons and 
one neutron, whereas an ordinary helium nucleus has two protons and 
two neutrons. This step of producing 3 He from a deuteron occurs very 
rapidly, so that a typical deuteron in the core of the Sun will survive for 
only 4 seconds before reacting with a proton. 

Usually, the final reaction in the proton-proton chain requires the addition 
of another neutron to the 3 He nuclei, thereby making normal 4 He. This 
final step can proceed in several ways, with the most common involving a 
collision of two 3 He nuclei. Each of these 3 He nuclei resulted from a prior, 
separate occurrence of step 2 somewhere else in the core. The final result is 
a normal 4 He nucleus and two protons. On average, a 3 He nucleus must wait 
4 million years before it participates in this reaction. 

The symbol v indicates a neutrino, the /3 + indicates a positron, and y indicates 
a photon. The orange sphere represents protons, and blue indicates neutrons. 

The net result of the chain of reaction is: 

6 1 H -» 4 He+2 1 H+26.72MeV 

Although it takes six protons to make one helium nucleus, there is a net loss of 
only four protons because two are regenerated in the final step. Because the six 
protons are more massive than the two protons and a helium nucleus, mass is lost 
in the proton-proton chain and converted to energy. Each resulting 4 He nucleus 
has a mass that is slightly less than the combined mass of the four protons that 
created it (by about 0.7%). The energy produced by a single proton-proton chain 
reaction is 26.27 MeV, and although the units are unfamiliar to you, this converts 
to about one ten-millionth of the amount of energy needed to lift a drop of water. 
As you can see, this is not a lot of energy; overall, however, the Sun converts 
about 600 million tons of hydrogen into 596 million tons of helium every second. 
The missing 4 million tons of matter are converted to energy in accordance with 
the famous equation formulated by Einstein: E = me 2 . The neutrinos carry off 
about 2% of this energy and rarely interact with matter and so pass straight out 
into space. The remaining energy emerges as kinetic energy of the nuclei and as 
radiative energy of the gamma rays. 


90 


Step 2: 


Step 3: 


Box 3.1: Mass and Energy Conversion 
in the Sun 

It is easy to calculate how much mass the Sun loses through nuclear fusion. First let 
us look at the input and output masses of the proton-proton chain A single proton 
has a mass of 1.6726 x l(U 27 kg, so four protons have a mass of 6.693 x 10 -27 kg. 



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A single 4 He nucleus has a mass of only 6.645 x 10 27 kg, which is slightly less than 
the mass of four protons. In other words: 

6.69 x 10' 27 kg-6.643 x 10~ 27 kg = 4.8 x 10" 29 kg 

This is only 0.7% or 0.007 of the original mass, so if 1 kg of hydrogen fuses, the 
resulting helium weighs 993 grams, and 7 grams of mass are turned into energy. To 
calculate the amount of energy for a single reaction, let us use Einstein’s famous 
equation, E = me 2 . 

E = me 2 = (4.8 x 10~ 29 kg) (3.0 x 10 8 ms -1 ) 2 = 4.3 x 10" 12 joules 

This is the tiny amount of energy created by the formation of one helium atom. It 
is so small, it would only power a 10-watt light bulb for half a trillionth of a second. 

Let us now see how much total energy is produced when hydrogen is converted to 
helium every second. We know that only 0.7% of the mass of hydrogen is fused: 

E = me 2 = (0.7kg) (3.0 x 10 8 ms -1 ) 2 = 6.3 x 10 14 joules 

However, the Sun’s luminosity is 3.9 10 26 J/s, thus hydrogen must be consumed at 
a rate: 

(3.9 x 10 26 joules per second) 

(6.3 x 10 14 joules per second) 

= 6x 10 11 kg/s 

So the Sun fuses 600 million metric tons of hydrogen each second, with 596 tons 
fused into helium and the remaining 4 million becoming energy. 


3.6.3 Energy Transport from the Core 
to the Surface 

Energy produced in the central region of the Sun flows outward toward the 
surface. If the Sun were transparent, the photons, or gamma rays, emitted by 
the extremely hot gases in the core would travel straight out at the speed of 
light, 2 seconds after being emitted. The Sun’s gases, however, are not very 
transparent, and so a typical photon only travels about 10 -6 meters before it is 
re-absorbed. In being absorbed, it heats up the surrounding gases, which in turn 
emit photons, which are then subsequently re-absorbed. The emitted photon will 
not necessarily be emitted in an outward direction but rather a totally random 
direction, which means that at least 10 25 absorptions and re-emissions occur 
before energy reaches the surface. This slow, outward migration of photons is 
often called a random walk. 

This process means that there is a considerable time delay before energy 
produced at the core reaches the surface. On an average, about 170,000 years 
will pass before energy created at the core eventually reaches the surface. 23 



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Furthermore, the energy produced in one second does not erupt from the surface 
all in one go; it appears that it is radiated from the surface over a period of more 
than 100,000 years. Some energy appears in about 120,000 years, while other 
energy takes 220,000 years. But the bulk of it is emitted after 170,000 years. 

This tells us two things about the Sun. First, when we observe the light emitted 
by the Sun, we learn nothing about what is going on in the core at that moment. 
All we can say is that energy was created in the core thousands of years ago. The 
second point is that if energy generation were to suddenly cease in the core for, 
say, a day, or even a hundred years, we would not notice it because by the time 
the energy flowed to the surface, it would have been averaged out over more 
than 100,000 years. This implies that the brightness of the Sun is very insensitive 
to changes in the energy production rate. 

The above processes occur in some form or other in many of the stars on the 
main sequence. As we shall see, more massive stars carry their energy outward 
in a different manner, and the energy is created in a slightly different way. We 
shall now look at other stars and how they are placed on the main sequence. 

Observing the Sun is a very popular pastime for amateur astronomers, but 
let me say now that you should NEVER OBSERVE, OR EVEN LOOK AT, THE 
SUN WITH THE NAKED EYE OR THROUGH A TELESCOPE. It is exceedingly 
dangerous, and you must have specially made equipment to do so. So do not 
do it. Instead, project the Sun onto a card. There are several excellent books on 
solar observing, and I highly recommend a very recent one by Chris Kitchin, 
called “Observing the Sun.” 


3.7 Binary Stars and Stellar Mass 


3.7.1 Binary Stars 

Binary stars, or as they are sometimes called, double stars, are stars that may 
appear to the naked eye to be just one star, but on observation with either binoc- 
ulars or telescopes resolve themselves into two stars. Indeed, some apparently 
single stars turn out to be several stars! Many appear as double stars due to their 
position in the same line of sight as seen from the Earth, and these are called 
optical doubles. It may well be that the two stars are separated in space by a vast 
distance. 

Others, however, are actually gravitationally bound and may orbit around each 
other over a period of days, or even years. These systems are the ones we will 
discuss here. 24 

The classification of some binary stars is quite complex. For instance, many 
cannot be resolved by even the largest telescopes and are called Spectroscopic 
Binaries, the double component only being fully understood when the spectra are 
analyzed. Others are Eclipsing Binaries, such as Algol ((3 Persei), where one star 
moves during its orbit in front of its companion, thus brightening and dimming 
the light observed. A third type is the Astrometric Binary, such as Sirius (ct Canis 
Majoris), where the companion star may only be detected by its influence on the 
motion of the primary star. As this book is concerned with objects that can be 



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observed visually, I will concentrate on binary stars that are physically associated 
and can be split with either the naked eye or with the use of some sort of optical 
equipment. 

What makes binary stars so important to astronomers is that by observing 
their motion, 25 as they dance around each other, it is possible to determine 
their mass. This is, of course, vitally important in determining the evolutionary 
processes of stars. 

Terminology must now be introduced that is specific to visual binary star 
observation. The brighter of the two stars is usually called the primary star, while 
the fainter is called the secondary (in some texts it may be called the companion, 
and both terms will be used throughout this book). This terminology is employed 
regardless of how massive either star is or whether the brighter is in fact the less 
luminous of the two but simply appears brighter. 

Perhaps the most important terms used in visual binary star work are the 
separation and position angle (PA). The separation is the angular distance 
between the two stars, usually in seconds of arc, and is measured from the brighter 
star to the fainter. The position angle is a somewhat more difficult concept to 
understand. It is the relative position of one star, usually the secondary, with 
respect to the primary, and it is measured in degrees, with 0° at due north, 90° at 
due east, 180° at due south, 270° at due west, and back to 0°. It is best described 
by an example: using Figure 3.9, the double star y Virginis, with components 
of magnitude 3.5 and 3.5, has a separation of 1.8" (arcsecs) at a PA of 267° 


2070 




94 


Astrophysics is Easy 


(epoch 2000.0). Note that the secondary star is the one always placed somewhere 
on the orbit, the primary star is at the center of the perpendicular lines, and 
the separation and PA of any double star are constantly changing and should 
be quoted for the year observed. Some stars (where the period is very long) will 
have no appreciable change in PA for several years; others, however, will change 
from year to year. 

It is worth mentioning again that although your optical equipment, including 
your eyes, should in theory be able to resolve many of the binaries listed here, 26 
there are several factors that will constrain the resolution (e.g., the seeing condi- 
tions, light pollution, dark adaption, your temperament, etc.). Thus, if you cannot 
initially resolve a double star, do not despair, but move onto another and return 
to the one in question at another date. Also recall that the colors ascribed to 
a star will not necessarily be the colors you see; they are just indicators of the 
general color, and in fact, as you will see from the text, many observers disagree 
on several stars’ colors. 

Presented below is a brief list of several stars that will, as an initial 
indicator, help you to determine the resolution of both yourself and your binoc- 
ulars/telescope. 27 All of the positions quoted are for the primary star. 

3.7. 1 . 1 Visual Binary Stars 

p, Canis Majoris ADS 5605 06 h 56.1' n — 14°03' January 4 
5.3,8.6m P.A. 340°; Sep. 3.0" Spec.G5 moderate 

Two stars of differing brightness that nevertheless present a glorious double of 
orange and blue. 

| Ursae Majoris ADS 8119 U h 18.2 m +31°32' March 12 

4.3,4.8m P.A. 273°; Sep. 1.8" Spec. GO very difficult 

Discovered by William Herschel in 1780, this is a close pair of pale yellow stars. It 
also has the distinction of being the first binary system to have its orbit calculated 
by Savary in 1828. Both components are also spectroscopic binaries. 

£ Ursae Majoris ADS 8891 13 ,1 23.9 m +54°56' April 12 

2.3,4.0m P.A. 152°; Sep. 14.4" Spec. A2 A2 very easy 

Part of the famous double Mizar and Alcor (80 UMa). Visible to the naked 
eye. Nice in binoculars. A small telescope will resolve Mizar’s 4th-magnitude 
companion. Alcor and both members of Mizar are themselves spectroscopic 
binaries. Thus, there are six stars in the system. Mizar also has several other 
distinctions: the first double to be discovered by telescope (by Riccioli in 1650), 
the first to be photographed (by Bond in 1857), and the first spectroscopic binary 
to be detected (by Pickering in 1889). 

a CVn ADS 8706 12 , '56.0 m +38°19' April 5 

2.9,5.5m P.A. 229°; Sep. 19.4" Spec. A0 easy 


Stars 


95 


Also known as Cor Caroli, the stars of this system are separated by a distance 
equivalent to 5 solar system widths — 770 astronomical units! The two stars are 
yellowish in small instruments; however, with large aperture, subtle tints become 
apparent and have been called flushed white and pale lilac or pale yellow and 
fawn! 


e Bootis ADS 9372 13 , '45.0 m +27°04' April 18 
2.5,4.9m P. A. 339°; Sep. 2.8" Spec. KO AO moderate 

Also known as Mirak. A wonderful contrast of gold and green stars has also been 
reported to be yellow and blue. Difficult with apertures of about 7.5 cm, and even 
a challenge for beginners with apertures of 15.0 cm. With small telescopes, a high 
power is needed to resolve them. 

(3 Lyrae ADS 11745 18'’50.1 m +33°22' July 4 

3.4 V , 8.6m P. A. 149°; Sep. 45.7" Spec. B9 easy 

This pair of white stars is a challenging double for binoculars. /3 1 is also an 
eclipsing binary. A fascinating situation occurs due to the gravitational effects 
of the components of (3 1 : the stars are distorted from their spherical shapes into 
ellipsoids. 


(3 Cygni ADS 12540 19 ,, 30.7 m +27°58' July 14 

3.1,5.1m P.A.54°; Sep. 34.4" Spec. K3 B8 easy 

Thought by many to be the finest double in the skies, Albireo is a golden-yellow 
primary and blue secondary against the backdrop of the myriad of fainter stars 
of the Milky Way. Easy to locate at the foot of the Northern Cross. 

e Lyrae ADS 11635 18 ,! 44.3 m -39° 40' July 2 

5.4,6.5m P. A. 357°; Sep. 2.6" Spec. A2 F4 easy/moderate 
5.1,5.3m P. A. 94°; Sep. 2.3" easy/moderate 

The famous Double-Double, easily split, but to resolve the components of each 
star, s 1 (magnitude 4.7) and e 1 (magnitude 4.6), requires a high power. The stars 
themselves are at a P.A. 173°, separated by 208", which is near the naked-eye 
limit, and some keen-eyed observers report being able to resolve them under 
perfect seeing conditions. However, there is fierce debate among amateurs — 
some say the double is difficult to resolve, others the opposite. All stars are white- 
or cream-white-colored. A highlight of the summer sky. 

61 Cygni ADS 14636 21 ,1 06.9 m -38° 45' August 8 
5.2,6.0m P.A. 150°; Sep. 30.3" Spec. K5 K7 easy 

Best seen with binoculars (but sometimes a challenge, if conditions are poor), which 
seem to emphasize the vibrant colors of these stars, both orange-red. Famous for 
being the first star to have its distance measured by the technique of parallax. The 
German astronomer Friedrich Bessel determined its distance to be 10.3 l.y.; modern 


96 Astrophysics is Easy 

measurements give a figure of 11.36. Also has an unseen third component, which 
has the mass of 8 Jupiters, and a very large proper motion. 

a Ursae Minoris ADS 1477 02 h 31.8 m +89°16' October 29 

2.0,8.2m P.A. 218°; Sep. 18.4" Spec. F8 easy 

Possibly the most famous star in the sky. Polaris, or the Pole Star, is located less 
than a degree from the celestial pole and is a nice double consisting of a yellowish 
primary and a faint whitish-blue secondary. The primary is also a Population II 
cepheid variable and a spectroscopic binary. Although claims have been made to 
the effect that the system can be resolved in an aperture as small as 4.0 cm, at 
least 6.0 cm will be required to split it clearly 

o 2 Eradini ADS 3093 04''15.2 m -07°39' November 24 

4.4,9.5m P.A. 104°; Sep. 83" Spec. WD easy/moderate 

A challenge to split with binoculars. What makes this system so interesting is 
that the secondary is the brightest white dwarf star visible from Earth. 

3.7.2 The Masses of Orbiting Stars 

It may come as no surprise to you that the mass of a star can be determined. 
However, the question that needs to be asked is “how”? Well, usually, we need 
to use binary stars, as well as the laws of Kepler and Newton. 

Kepler’s law, which demonstrates how the time required for a planet orbiting 
the Sun is related to its distance from the Sun, can be modified to describe the 
motion of any two bodies that orbit around each other. The person who first did 
this was the great Isaac Newton. To find the mass of the stars in a visual binary, we 
must first determine their orbits by observing them over several years. This might 
take a few years, or even tens of years, but eventually we can determine the time 
needed for one star to completely orbit the other. This period of time is called P. 
By using a plot of the orbit, and forearmed with the system’s distance from the 
Sun, we can then measure the semimajor axis, 28 a, of one star from the other. An 
important point to note here is that this method only gives us the combined mass of 
the stars, not their individual masses. To achieve that, we need to go one step further. 

From the above description, we can easily determine the combined masses of 
two stars orbiting each other. To determine individual stellar masses, however, 
we determine how much one star moves relative to the other. For instance, if 
one star is much more massive than the other, then it will hardly move at all 
relative to the less-massive star; the less-massive star will do all of the orbiting in 
the system in a manner reminiscent of planets orbiting the Sun. To be accurate, 
we should really say that the stars (and, incidentally, the planets and Sun) orbit 
about their common center of mass, or center of gravity. In fact, they “wobble.” 29 
The masses of stars are not so similar to that of, say, the Sun and Jupiter, 
where 99% of the mass is in the Sun, so they orbit the center of mass that is 
more or less equidistant from both of them. This center of mass is found along a 
line joining the two stars at a position that depends on the stars’ masses. Think 
of it as the balance point on a child’s see-saw. If one star is four times as massive 


Stars 


97 


as the other, the balance point will be four times closer to the more massive 
star. If the two stars have mass M A and M B , their distances from the center of 
mass (the balance point) are a A and a B , 30 with the larger mass having the smaller 
distance from the center of mass. So, for example, if the two stars have equal 
mass, or M A = M B , then a A — a B , and they orbit a point that is exactly halfway 
between them. On the other hand, if star B is four times less massive than star 
A(M b = 1/4 x M a ), star B orbits four times farther from the center of mass than 
star A (a B = 4 x a A ). Again, using the image of a see-saw, an adult weighing four 
times as much as a child must sit four times closer to the pivot as a child. 

In this manner, stellar masses can be determined, and using sophisticated 
techniques, the masses of double-star systems that cannot be optically resolved 
can also be measured. Using these and other techniques, we have determined 
that stellar mass ranges from about 0.08 M G to 50 M 0 . Of course, there are larger 
stars, but these are few and far between. 


Box 3.2: Determining Stellar Mass 

Consider the orbits of the double-star system of Sirius A and Sirius B. 

The two stars have an orbital period, P, of 50.1 years and an average semimajor 
axis, a, of 19.6 astronomical units. To determine their combined mass, M A +M B , we 
use the modified form of Kepler’s Law: 


M A+M B =y 2 

Inserting the measured values for P and a, we get: 


m a+ m b = 


19.8 3 

50. 1 2 


^7762 
~ 2510 
=3.1 M 0 . 


Thus, the combined mass of Sirius A and Sirius B is approximately 3.1 M 0 . In reality, 
the orbit of the system is very eccentric, or elliptical, and so the distance between them 
varies from 31.5 AU to 8.1 AU and back again. Using this information, and data gained 
from spectroscopic studies, we can determine that the mass of Sirius A is 2.12 M 0 , and 
that of Sirius B 1.03 M 0 . 


3.8 Lifetimes of Main-Sequence 
Stars 


We have covered topics that describe how a star forms, how the mass of stars 
can be determined by observing binary-star systems, and how long it takes to 
become a star. Now we shall discuss how long a star will remain on the main 
sequence, and then look at what happens due to changes in its internal structure. 



98 


Astrophysics is Easy 


The stars that are on the main sequence are fundamentally alike in their cores 
because it is here that stars convert hydrogen to helium. This process is called 
core hydrogen-burning. The main-sequence lifetime is the amount of time a star 
spends consuming hydrogen in its core, and so the main-sequence lifetime will 
depend on the star’s internal structure, and evolution. A newly born star is 
often referred to as a zero-age-main-sequence-star, or zams for short. There is a 
subtle but important difference between a zams star and a main-sequence star. 
During its long life on the main sequence, a star will undergo changes to its 
radius, surface temperature, and luminosity due to the core hydrogen-burning. 
The nuclear reactions alter the percentage of elements within the core. Initially, 
it would have had, say, in the case of the Sun, about 74% hydrogen, 25% helium, 
and 1% metals, but now, after a period of 4.6 million years, the core has a much 
greater mass of helium than that of hydrogen at its core. 

Due to the hydrogen-burning at the core, the total number of atomic nuclei 
decreases with time, and so with fewer particles in the core to provide the internal 
pressure, the core will shrink very slightly under the weight of the star’s outer 
layers. This has an effect on the star’s appearance. The outer layers expand and 
become brighter. This may seem odd to you; if the core shrinks, why doesn’t 
the star shrink? The explanation is very simple: the core shrinkage increases its 
density and temperature, which causes the hydrogen nuclei to collide with each 
other much more often, which in turn increases the rate of hydrogen-burning. 
The resulting increase of core pressure causes the star’s outer layers to expand 
slightly, and as luminosity is related to the surface area of a star, the increase 
of the star’s size will result in an increase in luminosity. In addition, the surface 
temperature will increase. In the case of the Sun, astronomers have calculated that 
its luminosity has increased by 40%, its radius by 6%, and its surface temperature 
by 300 K, all during the past 4.6 billion years. 

As a star ages on the main sequence, the increase of energy flowing from its 
core will also heat the surrounding area, and this will cause hydrogen-burning to 
begin in this surrounding layer. As this can be thought of as “new” fuel for the 
star, its lifetime can be lengthened by a few million years for a main-sequence star. 

The one factor that determines how long a star will remain on the main 
sequence is its mass. Basically, it can be summed up in a few words: Low-mass 
stars have much longer lifetimes than high-mass stars. Figure 3.10 illustrates this 
nicely. 

High-mass stars are extremely bright, and their lifetimes are very short. This 
means that they are using up their reserve of hydrogen in the core at a very 
high rate. Thus, even though an O- or B-type star is much more massive, and 
thus contains more hydrogen than, say, a less-massive M-type star, it will use 
up its hydrogen much sooner. It may only take a few million years for O- 
or B-type stars to use up their supply of hydrogen, whereas for low-mass M- 
type stars, it may take hundreds of billions of years. Think about that for a 
second. The lifetime of an M-type star may be longer than the present age of the 
universe ! 31 Table 3.1 shows how the mass of a star relates to temperature and 
spectral class. 

The differing lifetimes of stars can be easily seen by looking at star clusters. 
Massive stars have shorter lifetimes than less-massive stars, and so a star 
cluster’s H-R diagram will give information on the evolution of the stars in 


Stars 


99 



Mass (solar masses) 


Figure 3. TO. Main-sequence lifetimes for stars of different mass. 


the cluster. Such a diagram will show a main sequence that lacks O-type stars, 
which are the most massive, then A-type, and so on and so forth as the 
cluster ages. Figure 3.11 shows this erosion of main-sequence stars by comparing 
the H-R diagrams of several different star clusters. In every case, some stars 
will have left the main sequence to become red giants, 32 or they are already 
red giants. 

The temperature and spectral type of the very hot stars that are left on the 
main sequence are used to determine the age of a star cluster. Suppose the 
hottest star on the main sequence is an AO-type star, with the much hotter, and 
more-massive stars already evolved to red giants. We know that AO stars have 
a main-sequence lifetime of about 100 million years, so we can say with some 
confidence that the star cluster is about 100 million years old. 

Generally, the more massive the star, the faster it goes through all of its phases, 
so we are fortunate to be able to observe stars in the main-sequence phase, as 
they remain in it for such a long time. It is also very easy to estimate a star’s 
lifetime if we know its mass. 


Table 3.1 

. Mass, spectral class, and main-sequence lifetimes 


Mass, 

Temperature, 

Spectral Class 

Luminosity, 

Main-Sequence Lifetime, 

m g 

K 


Lg 

1 0 6 years 

25 

35,000 

O 

80,000 

3 

15 

30,000 

B 

10,000 

1 1 

3 

1 1 ,000 

A 

60 

640 

1.5 

7000 

F 

5 

3600 

1 

6000 

G 

1 

10,000 

0.75 

5000 

K 

0.5 

20,000 

0.5 

4000 

M 

0.03 

56,000 



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Astrophysics is Easy 


Luminosity (L 0 ) 



Temperature (K) 


Figure 3.11. H-R diagrams for several star clusters of different ages. 


There are many stars on the main sequence that can be observed. The brightest 
of these have already been mentioned in previous sections; they include: Regulas, 
Vega, Sirius A, Procyon A, the Sun, and Barnard’s Star, to name a few. 


Box 3.3: Main-Sequence Lifetimes 

The duration a star remains on the main sequence is very easy to calculate. There is 
an approximate relationship between the mass of a star and its lifetime: 

1 _ 1 

“ M 2 - 5 “ M 2 VM 


Astronomers usually relate the main-sequence lifetime to the Sun (a typical 1 M 0 
star), which is believed to be 10 10 years, or ten billion years. 

For example, the main-sequence lifetime of Sirius, a 2.12M 0 , star will be: 


1 

2 . 12 2 - 5 


-= = solar lifetimes 

2.12 2 v / 2d2 6.54 




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101 


So, a Sirius will burn hydrogen in its core for about 1/6.54 x 10 10 years, or about 
1.5 billion years. 

On the other hand, a main-sequence star with a mass of 0.5 M e will have a lifetime 
of: 


1 

05 “ 


r= = = 5.66 solar lifetimes 

0.5 2 -s/(h5 0.177 


which is about 56 billion years. 


3.9 Red Giant Stars 


Although the amount of hydrogen in a star’s core is vast, it is not infinite, and 
so, after a very long time , the production of energy will cease when the central 
supply of hydrogen is used up. Throughout the length of time that nuclear fusion 
has been taking place, the hydrogen has been transformed into helium, by way 
of the proton-proton chain, and without this source of energy, the star uses 
gravitational contraction to supply its energy needs. Thus, the core will start to 
cool down, which means that the pressure also decreases, with the result that the 
outer layers of the star begin to weigh down on the core and compress it. This 
has the effect of causing the temperature within the core to rise again, and for 
heat to flow outward from the core. Note that although a tremendous amount of 
heat is formed now, it is not due to nuclear reactions but to gravitational energy 
being converted into thermal energy. 

In a relatively short time, astronomically speaking, the region around the 
star’s hydrogen-depleted core will become hot enough to begin nuclear fusion of 
hydrogen into helium, in a thin shell around the core, in a process called shell 
hydrogen-burning. This is shown in Figure 3.12. 

The core will consist of helium, but the outer layers are hydrogen rich. The 
shell, where energy production occurs, is relatively thin. [Note: figure to scale.] 

For a star like the Sun, this hydrogen-consuming shell develops almost immedi- 
ately from the moment nuclear fusion stops in the core, and so the supply of 
energy is more or less constant. For massive stars, there can be an interval of 
perhaps a few thousand years to a few million years from the end of the core 
nuclear-fusion phase to the beginning of the shell hydrogen-burning phase. 

The new supply of energy, and thus heat, has the effect of causing the rate 
of shell hydrogen-burning to increase, and so it begins to eat further into the 
surrounding hydrogen. The helium that is the by-product of the hydrogen fusion 
in the shell falls to the center of the star, where, along with the helium already 
there, it heats up as the core continues to contract and increase its mass. In the 
case of, say, a 1 M 0 star, the core will be compressed to as much as one-third 
of its original size. The result of this core compression is an increase in the 
temperature, from about 15 million K to nearly 100 million K. 

Now, most of what has happened in this stage of a star’s life has occurred 
inside of it and so was invisible to our eyes. Nevertheless, it does have effects 





102 


Astrophysics is Easy 


Hydrogen-burning shell 



Hydrogen-rich outer layers 


Figure 3.12. Star with shell hydrogen -burning. 


on a star’s structure that can drastically alter its appearance. The star’s outer 
layers expand as the core contracts. With the increased flow of heat from 
the contracting core, and the ever-expanding shell of hydrogen-burning, the 
star’s luminosity increases quite substantially. This causes the star’s internal 
pressure to increase and makes the outer layers of the star expand to many 
times their original radius. The tremendous expansion actually causes the outer 
layers to cool, even though the inner core temperature has risen dramati- 
cally. The new, much-expanded, and cooler outer layers can reach tempera- 
tures as low as 3500 K and will glow with a very distinctive reddish tint, as 
can be explained by Wien’s Law (mentioned earlier). The star has now become 
a red giant star. 

So, we can now see that red giant stars are former main-sequence stars that 
have evolved into a new phase of their lives. 

Due to the large diameter and thus weaker surface gravity of the red giant, quite 
a substantial amount of mass loss can occur. This means that gases can escape 
from the surface of a red giant star. Such an effect is relatively easy to observe 
by looking at the absorption lines produced in the star’s spectrum. Calculations 
and measurements have shown that a typical red giant star can lose 10 -7 M o 
each year. Compare this with the much-smaller 10“ I7 M o that the Sun loses each 
year. From this, we can see that as a star evolves from the main sequence to the 
red-giant stage, it can lose quite a lot of its mass. The evolutionary track from 
the main sequence to the red-giant phase for stars of differing mass is shown in 
Figure 3.13. 

The dotted lines indicate time scales of 10, 50, 100 million, and 1 billion 
years. You can see that a star of about 15 solar masses leaves the main 
sequence (the shaded area) about 100 times earlier than a star of 1.5 solar 
masses. 

There are many wonderful red giant stars that are observable in the night 
sky. We have already mentioned a few of these in earlier sections — Capella A, 



Stars 


103 



Temperature (K) 

Figure 3.13. Evolutionary tracks from the main sequence to the red-giant phase for stars of 
different mass. 


Arcturus, Aldebaran, Pollux, Mirach, R Leporis, and Mira. But there are several 
other, lesser-known red giant stars, which are also worth observing. 


3.9.1 Bright Red Giant Stars 

RS Cyg HD 192443 20 h 13.3 m +38° 44' Jun-Jul-Aug 

8.1 v m B-V:3.3 C5 Cygnus 

A red giant star with a persistent periodicity, class SRA, it has a period of 
417.39 days, with a magnitude range of 6.5 to 9.5 m. A strange star where the 
light curve can vary appreciably, with the maxima sometimes doubling a deep 
red-colored star. 

R Aqr HD 222800 23' ! 43.8 m -15° 17' Jul-Aug-Sep 

5.8 v m B-V:1.5 M4 pe 2500 Aquarius 


This is a symbiotic double star and is classed as a Z Andromedae - type star. 
R Aqr is a nice red giant, which, incidentally, has a small, blue (thus, very hot) 



104 


Astrophysics is Easy 


companion star. Due to its variable nature, its magnitude can fall to 11.5, and so 
it can be somewhat difficult to locate. It is believed to lie at a distance of about 
640 l.y. 


R Cas HD 224490 23''58.4 m +51°23' Sep-Oct-Nov 

5.5 v m B-V:1.5 M7 Hie 2000 Cassiopeia 

This is a Mira-type variable star with quite a large magnitude range, say 
5.5-13.0m. It is estimated to lie at a distance of 350 l.y. Its surface temperature 
of only 2000 K is still a matter of speculation. 

R Leo HD 84748 09*47.6” +11°26' Jan-Feb-Mar 

6.02 v m B-V : 1.5 M8 Hie 2000 Leo 

A very bright Mira-type variable star and a favorite among amateur astronomers. 
Again, like so many other red giants, its low temperature of 200 K is in some 
doubt. Its color is deep red. This star is often cited as being a perfect introductory 
star for those who wish to observe a variable star. It is also an AGB star. 

Note that as a star ages and moves from the main sequence to the red-giant 
stage, its spectral type will also change. For instance, the Sun, at present a 
G-type star, will gradually change its spectral class to K, and then to a warm 
M-type star. It may even become an M2- or M3-type with the temperature 
falling to about 3200 K. Similarly, stars of different masses will also change their 
spectral type. 

Finally, for this section, it is worth noting in passing that there are two red 
giant phases. Which one a star will follow depends, as I am sure you will have 
guessed by now, on its mass, and can lead to the formation of supergiant stars. 
We will discuss these in another section. 


3.10 Helium-Burning 

and the Helium Flash 


Most of the stars with a mass greater than, or equal to, the Sun’s will eventually 
become red giants. But how energy is produced in the star after it has reached the 
red-giant phase depends on its mass. We shall look at these two stages, beginning 
with how the helium in its core produces energy. 

3.10.1 Helium-Burning 

Helium can be thought of as the “ash” left over from the hydrogen-burning 
reactions, and can in fact be used as the fuel for another nuclear fusion reaction, 
which, this time, uses helium. This is the helium-burning phase. As a star 
approaches and becomes a red giant, its core temperature is too low to initiate 
helium-burning. But the hydrogen-burning shell that surrounds the dormant 
helium core adds mass to the core, with the result that it contracts further, 


Stars 


105 


becomes denser, and increases the temperature substantially. 33 Something else 
happens as the temperature increases — the electrons in the gas become degen- 
erate. Electron degeneracy is a very important process and is explained in greater 
detail in the appendices. When the electrons become degenerate, they in effect 
resist any further contraction of the core, and the internal temperature of the 
core will no longer affect the internal pressure. 

As the hydrogen shell continues to burn, the degenerate core grows even hotter, 
and when it reaches 100 million K, and has a mass of about 0.6 M 0 (i.e., the 
inner 60% of the hydrogen in the star has been converted to helium), core 
helium-burning begins, converting helium into carbon, and producing nuclear 
energy. During this stage of a star’s life, it can be nearly 1 AU in radius and 
almost 1000 times as luminous as the Sun. By now the old star has once again 
obtained a central energy source for the first time since it left the main sequence. 

The helium-burning in the core fuses three helium nuclei to form a carbon 
nuclei and is called the triple a process. This occurs in two steps. In the first step, 
two helium nuclei combine to form an isotope of beryllium: 

4 He+ 4 He -» 8 Be 

This isotope of beryllium is very unstable and very quickly breaks into two helium 
nuclei. But in the extreme conditions in the core, a third helium nucleus may 
strike the 8 Be nucleus before it has had a chance to break up. If this happens, 
a stable isotope of carbon is formed, and energy is released as a gamma-ray 
photon (y): 

8 Be+ 4 He -* 12 C+y 

The phrase “triple a” comes about because helium nuclei are also called alpha 
particles. 34 The carbon nuclei formed in this process can also fuse with additional 
helium nuclei, producing a stable isotope of oxygen, and supply additional energy: 

12 C+ 4 He -» 16 0+y 

So the “ash” of helium-burning is carbon and oxygen. This process is very 
interesting, as you will note that both of these isotopes of oxygen and carbon 
are the most abundant forms and in fact make up the majority of carbon atoms 
in our bodies, as well as the oxygen we breathe. We will explore this fascinating 
piece of information in greater depth later in this book. 

The formation of carbon and oxygen not only provides more energy but also 
re-establishes thermal equilibrium in the core of the star. This prevents the core 
from any further contraction due to gravity. The duration of time a red giant 
will spend burning helium in its core is about 20% as long as the time it spent 
burning hydrogen on the main sequence. The Sun, for example, will only spend 
2 billion years in the helium-burning phase. 

3.10.2 The Helium Flash 

As I mentioned earlier, the mass of a star will direct how helium-burning begins 
in a red giant star. In a high-mass star (that is, with a mass greater than 


106 


Astrophysics is Easy 


2-3 M 0 ), the helium-burning begins gradually as the temperature in the core 
approaches 100 million K. The triple a process is initiated, but it occurs before 
the electrons become degenerate. However, in low-mass stars (that is, with a mass 
less than 2-3 M Q ), the helium-burning stage can begin suddenly, in a process 
called the helium flash. This stage, the helium flash, occurs due to the most 
unusual conditions found in the core of a low-mass star as it becomes a red 
giant. 

The energy produced by helium-burning heats up the core of the star and 
raises its temperature. Now, in normal circumstances, this would result in an 
increase of pressure that would lead to an expansion and subsequent cooling 
of the core. This explains why nuclear reactions do not usually cause a rapid 
increase in the central temperature of a star. But we must remember that the 
gas in the core of a 1 M Q red giant is far from normal; it is a gas of degenerate 
electrons. This means that any temperature increase that the helium-burning 
produces does not increase the internal pressure. What the rise in temper- 
ature does is to strongly affect the rate at which the triple a process occurs. A 
doubling of the temperature will increase the triple a production rate by about 1 
billion times. 

The energy that is produced by the triple a process heats up the core, and its 
temperature begins to rise even more. This increase and the subsequent rise in 
energy production can cause the temperature to reach an amazing 300 million K. 
Due to the rapid heating of the core, a nearly explosive consumption of helium 
occurs, and this is the helium flash mentioned earlier. At the peak of the helium 
flash, the core of the star has, very briefly, an energy output that is some 10 u 
to 10 14 times solar luminosity. This converts to a rate of energy output that is 
about 100 times greater than the entire Milky Way. 

Eventually, however, the high temperature becomes so high that the electrons 
in the core can no longer remain degenerate. They then behave normally for 
electrons in a gas, with the result that the star’s core expands, which ends the 
helium flash. These events occur very quickly, so the helium flash is over in a 
matter of seconds, and the star’s core settles down to a steady rate of helium- 
burning. 

An important point to make here is that whether the helium flash occurs 
or does not occur, the start of helium-burning actually reduces the star’s 
luminosity. Here’s what happens: the superheated core expands, and this core 
expansion pushes the hydrogen-burning shell outward, lowering its temperature 
and burning rate. The result is that even though the star has both helium fusion 
in its core and a shell of hydrogen-burning taking place simultaneously, the 
total energy production falls from its peak during the red-giant phase. This 
reduced total energy output of the star therefore reduces the luminosity and 
allows its outer layers to contract from their peak size during the red-giant 
phase. As the outer layers contract, the star’s surface temperature will increase 
slightly. 

The helium-burning in the core lasts for a relatively short time, however, and 
from calculations we can make an estimate of this time. For, say, a 1 M Q star like 
the Sun, the period after the helium flash will only last about 100 million years, 
which is 1% of its main-sequence lifetime. 


Stars 


3.1 1 Star Clusters, Red Giants, 
and the H-R Diagram 


At this point in our story of stellar evolution, it is a good idea to take stock of 
what we have learned so far. We have discussed how stars are formed before 
moving onto the main sequence. Their lifetime on the main sequence depends 
on their mass; massive stars have shorter lives. The red-giant phase is next, along 
with a change in the hydrogen, and helium-burning within the star’s core. To 
put all of this together in one coherent picture is useful, as we can see how a star 
develops from the moment of its birth; so we shall do just that by looking at the 
H-R diagram for stars that have just started their main-sequence lifetimes and 
those that are in the red-giant phase. 

Stars that have just emerged from the protostar stage and are about to join 
the main sequence are burning hydrogen steadily and have attained hydrostatic 
equilibrium. These stars are often referred to as zero-age main-sequence stars 
and lie along a line on the H-R diagram called the zero-age main sequence, 
or ZAMS. This is shown on the H-R diagram in Figure 3.14 as a green line. 
Over time, which can be relatively short or exceptionally long, depending on the 
star’s mass, the hydrogen in the core is converted to helium, and the luminosity 
increases. This is accompanied by an increase in the star’s diameter, and so the 
star moves on the H-R diagram away from the ZAMS. This explains why the 
main sequence is actually more of a broad band, rather than, as often portrayed, 
a thin line. 

The light grey line in Figure 3.14 represents those stars in which the hydrogen 
has been used up in the core and so nuclear fusion has ceased. As you can 
see, high-mass stars, 3M 0 , 5M 0 , and 10M o , then move rapidly from left 
(high temperature) to right (low temperature) across the H-R diagram. What 
is happening here is a decrease in surface temperature, but the surface area is 
increasing, so its overall luminosity remains fairly constant (i.e., an approximately 
horizontal line). In this phase, the core is contracting and outer layers expanding 
as energy flows from the hydrogen-burning shell. 

High-mass stars with core helium-burning exhibit sharp downward turns in 
the red-giant region of the H-R diagram. Low-mass stars have a helium flash at 
their cores (red stars). 

The evolutionary track of the high-mass stars then makes an upward turn to 
the upper-right section of the H-R diagram. This occurs just before the onset 
of core helium-burning. After the start of helium-burning, the core expands, the 
outer layers contract, and the evolutionary track of the star falls from these high, 
albeit temporary, luminosities. Notice how the tracks wander back and forth on 
the H-R diagram. This represents the stars’ adjusting to their new energy supplies. 

The low-mass stars, 1 M Q and 1.5 M 0 , behave in a somewhat different manner. 
The start of helium-burning is marked by the helium flash, indicated by the red 
stars in the diagram. The star shrinks and becomes less luminous after the 
helium flash, although the surface temperatures rises. This occurs because the 
reduction in luminosity is proportionally less than the reduction in size. So now 


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30,000 10,000 4,000 

Temperature (K) 


Figure 3.14. Post-main-sequence evolutionary track for several stars of different masses. 


the evolutionary tracks move down (lower luminosity) the H-R diagram, and 
to the left (hotter) region. 

We can observe the evolution of stars from birth to helium -burning by looking 
at young star clusters and comparing the actual observations with theoretical 
calculations. But there exists another astronomical grouping of stars that contain 
many, maybe millions, of very old post-main-sequence stars — globular clusters. 
These are the subject of our next section. 


3.12 Post-Main-Sequence Star 
Clusters: The Globular 
Clusters 


In the night sky there are many compact and spherical collections of stars. These 
stars clusters are called globular clusters. These are metal-poor stars and are usually 
to be found in a spherical distribution around the galactic center at a radius of about 
200 l.y. Furthermore, the number of globular clusters increases significantly the 
closer one gets to the galactic cente. This means that particular constellations which 
are located in a direction toward the galactic bulge have a high concentration of 
globular clusters within them, such as Sagittarius and Scorpius. 




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109 


The origin and evolution of a globular cluster is very different from that of an 
open or galactic cluster. All the stars in a globular cluster are very old, with the 
result that any star earlier than a G- or F-type star will have already left the main 
sequence and be moving toward the red-giant stage of its life. In fact, new star 
formation no longer takes place within any globular clusters in our Galaxy, and 
they are believed to be our Galaxy’s oldest structures. In fact, the youngest of 
the globular clusters is still far older than the oldest open cluster. The origin of 
globular clusters is a scene of fierce debate and research, with the current models 
predicting that they may have been formed within the proto-galaxy clouds that 
went to make up our Galaxy. 

As previously mentioned, globular clusters are old, as they contain no high- 
mass main-sequence stars, and this can be shown on a special kind of H-R 
diagram called a color-magnitude diagram. On a color-magnitude diagram, the 
apparent brightness is plotted against the color ratio for many of the stars in 
a cluster (see Figure 3.15). The color ratio of a star can tell you the surface 
temperature, and if we assume that all the stars in a cluster lie at the same 
distance from us, their relative brightness can tell us their relative luminosities. 

Even a cursory glance at such a color-magnitude diagram will tell you 
something strange has happened. In fact, you will see that the upper half of the 
main sequence has disappeared. This means that all of the high-mass stars in a 
globular cluster have evolved into red giants, a long time ago. What remains are 
the low-mass main-sequence stars that are very slowly turning into red giants. 

One thing that is very apparent on the diagram is a grouping of stars that lie 
on a horizontal band toward the center-left of the diagram. This is called the 
horizontal branch, while the stars in this area are the horizontal-branch stars. 
These stars are low-mass, post-helium-flash stars of about 50 L e , in which there 
are both core helium-burning and shell hydrogen-burning. In the future, these 
stars will move back toward the red-giant region as the fuel is devoured. 

A star that has had its surface temperature and visual magnitude determined 
is represented by a black dot. All the stars in M3 lie at approximately the same 
distance from us (32,000 l.y.), so their magnitudes are a direct measurement of 
their luminosities. The asymptotic giant branch is described in a later section. 

One very practical use of the H-R diagram is to estimate the age of a star 
cluster. With a very young star cluster, most, if not all, of the stars are on or near 
the main sequence. As it ages, however, the stars will move away from the main 
sequence, with the high-mass, high-luminosity stars being the first to become 
red-giant stars. As time passes, the main sequence will get increasingly shorter. 
The top of the main sequence, which remains after the specified time, can be 
used to determine the cluster’s age and is called the turnoff point. The stars that 
are at the turnoff point are those that are just exhausting the hydrogen in their 
cores, so the main-sequence lifetime is in fact the age of the star cluster. An 
example of the H-R diagram for open star clusters showing their turnoff points 
is shown in Figure 3.16. 

The time for the turnoff point is shown as a horizontal gray line. For example, 
the cluster M41 has a turnoff point near the 10 8 year point, so the cluster is about 
100,000,000 years old. 

From an observational point of view, globular clusters can be a challenge. 
Many are visible in optical instruments, from binoculars to telescopes, and a 


no 


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Absolute 

magnitude 

(M) 


Color ratio h v /h B 

0.75 1.0 1.5 2.0 2.5 3.0 4.0 



Surface temperature (K) 


Figure 3.1 5. Color-magnitude diagram for the globular cluster M3. 


few are even visible to the naked eye. There are about 150 globular clusters, 
ranging in size from 60 to 150 l.y. in diameter. They all lie at vast distances from 
the Sun and are about 60,000 l.y. from the Galactic plane. The nearest globular 
clusters (for example, Caldwell 86 in Ara) lie at a distance of over 6000 l.y., 
and thus the clusters are difficult objects for small telescopes. They cannot be 
seen; rather, any structure within the cluster will be difficult to observe. Even 
the brightest and biggest globular will require apertures of at least 15 cm for 
individual stars to be resolved. However, if large-aperture telescopes are used, 
these objects are magnificent. Some globular clusters have dense concentrations 
toward their center, while others may appear as rather compact open clusters. 
In some cases, it is difficult to say where the globular cluster peters out and the 
background stars begin. 



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Temperature (K) 

Figure 3.16. H-R diagram for star clusters showing the turnoff point. 


As in the case of open clusters, there exists a classification system, the Shapley- 
Sawyer Concentration Class, where Class I globular clusters are the most star- 
dense and Class XII the least. The ability of an amateur to resolve the stars in a 
globular actually depends on how condensed the cluster is, and so the scheme 
will be used in the descriptions, but it is really useful only for those amateurs who 
have large-aperture instruments. Nevertheless, the observation of these clusters, 
which are among the oldest objects visible to amateurs, can provide you with 
breathtaking, almost three-dimensional aspects. 

The many globular clusters listed below are just a few of the literally hundreds 
that can be observed and are meant to be just a representative sample. The 
© indicates the approximate size of the cluster. 


3. 1 2. 1 Bright Globular Clusters 

Messier 68 NGC 4590 12 h 39.5 m -26°45' Mar-Apr-May 

7.7m © 12' X Hydra 

Appearing only as a small, hazy patch in binoculars, this is a nice cluster in 
telescopes, with an uneven core and faint halo. A definite challenge to naked-eye 



Astrophysics is Easy 


112 

observers, where perfect seeing conditions will be needed. Use averted vision and 
make sure that your eyes are well and truly dark-adapted. 

Messier 3 NGC 5272 13 h 42.2 m +28°23' Mar-Apr-May 

5.9m © 16' VI Canes Venatici 

A good test for the naked eye. If using giant binoculars with perfect seeing 
conditions, some stars may be resolved. A beautiful and stunning cluster in 
telescopes, it easily rivals Ml 3 in Hercules. It definitely shows pale colored tints, 
and reported colors include yellow, blue, and even green; in fact, it is often 
quoted as the most colorful globular in the northern sky. Full of structure and 
detail, including several dark and mysterious tiny patches. Many of the stars in 
the cluster are also variable. One of the three brightest clusters in the northern 
hemisphere is located at a distance of about 35,000 l.y. 

Messier 5 NGC 5904 15 h 18.6 m +02°05' Apr-May-Jun 

5.7m © 17.4' V Serpens 

Easily seen as a disc with binoculars and with large telescopes, the view is 
breathtaking — presenting an almost three-dimensional vista. One of the few 
colored globulars, with a faint, pale yellow outer region surrounding a blue-tinted 
interior. It gets even better with higher magnification, and stars become more 
apparent. Possibly containing over half a million stars, this is one of the finest 
clusters in the northern hemisphere; many say it is the finest. 

Messier 4 NGC 6121 16 h 23.6 m -26°32' Apr-May-Jun 

5.8m © 26.3' IX Scorpius 

A superb object, presenting a spectacle in all optical instruments, and even visible 
to the naked eye. But it does lie very close to the star Antares, so the glare 
of the latter may prove a problem in its detection. Telescopes of all apertures 
show detail and structure within the cluster, and the use of high magnification 
will prove beneficial; but what is more noticeable is the bright lane of stars that 
runs through the cluster’s center. The closest globular to the Earth at 6,500 l.y. 
(although NGC 6397 in Ara may be closer), and about 10 billion years old. 

Messier 13 NGC 6205 16 h 41. 7 m +36°28' May-Jun-Jul 

5.7m © 16.5' V Hercules 

Also known as the Hercules Cluster. A splendid object and the premier cluster 
of the northern hemisphere. Visible to the naked eye, it has a hazy appearance 
in binoculars; with telescopes, however, it is magnificent, with a dense core 
surrounded by a sphere of a diamond-dust-like array of stars. In larger telescopes, 
several dark bands can be seen bisecting the cluster. It appears bright because it 
is close to us, at only 23,000 l.y., and also because it is inherently bright, shining 
at a luminosity equivalent to more than 250,000 Suns. At only 140 l.y. in diameter, 
the stars must be very crowded, with several stars per cubic light year, a density 
500 times that of our vicinity. 

Messier 10 NGC 6254 16 h 57.1 m -04°06' May-Jun-Jul 

6.6m © 15' VII Ophiucus 


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113 


Similar to M12, but slightly brighter and more concentrated. It lies close to 
the orange star 30 Ophiuchi (spectral type K4, magnitude 5), and so if you 
locate this star, then by using averted vision M10 should be easily seen. Under 
medium aperture and magnification, several colored components have been 
reported: a pale blue-tinted outer region surrounding a very faint pink area, with 
a yellow star at the cluster’s center. 

Messier 19 NGC 6273 17 h 02.6 m -26°16' May-Jun-Jul 

6.7m © 13.5' VIII Ophiucus 

A splendid, albeit faint, cluster when viewed through a telescope. Although a 
challenge to resolve, it is nevertheless a colorful object, reported as having both 
faint orange and blue stars, while the overall color of the cluster is a creamy 
white. 

Messier 9 NGC 6333 17 h 19.2 m -18°31' May-Jun-Jul 
7.6m © 9.3' VII Ophiucus 

Visible in binoculars, this is a small cluster with a brighter core. The cluster is 
one of the nearest to the center of our Galaxy and is in a region conspicuous 
for its dark nebulae, including Barnard 64; it may be that the entire region is 
swathed in interstellar dust, which gives rise to the cluster’s dim appearance. It 
lies about 19,000 l.y. away. 

Messier 22 NGC 6656 18 h 36.4 m -23°54' May-Jun-Jul 

5.1m © 24' VII Sagittarius 

A truly spectacular globular cluster, visible under perfect conditions to the naked 
eye. Low-power eyepieces will show a hazy spot of light, while high power will 
resolve a few stars. Often passed over by northern hemisphere observers due to 
its low declination. Only 10,000 l.y. away, nearly twice as close as M13. 

Messier 92 NGC 6341 17 h 17.1 m +43°08' May-Jun-Jul 

6.4m © 11' IV Hercules 

A beautiful cluster, often overshadowed by its more illustrious neighbor, M13. In 
binoculars, it will appear as a small hazy patch, but in 20 cm telescopes its true 
beauty becomes apparent, with a bright, strongly concentrated core. It also has 
several very distinct dark lanes running across the face of the cluster. A very old 
cluster, 25,000 l.y. distant. 

Messier 54 NGC 6715 18 h 55.1 m -30°29' Jun-Jul-Aug 

7.6m © 9.1' III Sagittarius 

It has a colorful aspect — a pale blue outer region and pale yellow inner core. 
Recent research has found that the cluster was originally related to the Sagittarius 
Dwarf Galaxy, but the gravitational attraction of our Galaxy has pulled the 
globular from its parent. Among the globular clusters in the Messier catalogue, 
it is one of the densest, as well as the most distant. 

Messier 15 NGC 7078 21 h 30.0 m +12°10' Jul-Aug-Sep 

6.4m © 12' IV Pegasus 


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An impressive cluster in telescopes, it can be glimpsed with the naked eye. It 
does, however, under medium magnification and aperture, show considerable 
detail, such as dark lanes, arcs of stars, and a noticeable asymmetry. It is one 
of the few globulars that have a planetary nebula located within it — Pease-1, 
which is seen only in apertures of 30 cm and greater. The cluster is also an X-ray 
source. 


3.13 Pulsating Stars 


We saw earlier that there are stars far more massive than the Sun that contract, 
and move horizontally across the H-R diagram, while at the same time they get 
hotter but remain at a constant luminosity. As they move across the H-R diagram, 
they can also become unstable and vary in size. Some stars change their size 
quite considerably, alternatively shrinking and expanding as their surface moves 
in and out. As the stars vary in size, so does their brightness. These stars are 
the pulsating variable stars. There exist several classes of pulsating variable stars, 
but we will just discuss the main types: the long-period variables, the Cepheid 
variables, and the RR Lyrae stars. Figure 3.17 shows where on the H-R diagram 
these pulsating stars reside. 



Figure 3.17. Variable stars on the H-R diagram. 




Stars 


115 


3. 1 3. 1 Why Do Stars Pulsate? 

You may think that the pulsations of a star are caused by variations in the rate of 
energy production deep in its core. You would be wrong, however, as the rate of 
nuclear fusion remains constant in a pulsating star. Astronomers have realized 
that the variations are caused by changes in the rate at which energy can escape 
from the star. The explanation is surprisingly simple, but somewhat involved, so 
I shall go through the various stages in some detail. 

Imagine a normal star, where there is a balance between the downward-pulling 
force of gravity and the upward force of pressure (i.e., the star is in hydrostatic 
equilibrium). Now picture a star where the pressure in the outer layers exceeds 
the force of gravity in those layers. In such a scenario, the star’s outer layers 
would begin to expand (see Figure 3.18 for a schematic of this process). As the 
star expands, its gravity will naturally fall, but the pressure force will fall at a 
faster rate. A time would then come when the star will have expanded to a larger 
size where, once again, hydrostatic equilibrium would reign. But this does not 
necessarily mean that the star would stop expanding. The inertia of the outward- 
moving outer layers will carry the expansion past the balance point. By the time 
gravity will have brought everything to a stop, the pressure would now be too 
small to balance the gravity, and so the outer layers would begin to fall inward. 
At this point gravity will rise again, but less than the pressure will. The outer 
layers will fall past the balance point until eventually the force of pressure would 
prevent any further fall, and so would come to a halt. And this is where we came 
in — the pulsations would start all over again. 

You can think of a pulsating star behaving just like a spring with a heavy 
weight attached to it. If you pull down on the weight and then let it go, the 
spring will oscillate around the point at which the tension in the spring and the 
force of gravity are in balance. After a while, however, friction in the spring will 
dampen the oscillations unless the spring is given a little push upwards each 
time it reaches the bottom of an oscillation. In a pulsating star, for the pulsations 
to continue, and not to die out, the star also needs an outward push each time 
it contracts to its minimum size. Discovering what causes that extra push was a 
challenge to astronomers of the twentieth century. 

The first person to develop an idea of what was happening was the British 
astronomer Arthur Eddington in 1914. He suggested that a star (in this case, a 
Cepheid variable) pulsated because its opacity increases more when the gas is 
compressed than when it is expanded. Heat is trapped in the outer layers if a 
star is compressed, which increases the internal pressure; this, in turn, pushes 
upward the outer layers. As the star expands, the heat will escape and so the 
internal pressure falls, and the star’s surface drops inward. 

In 1960, the American astronomer John Cox further developed the idea and 
proved that helium is the key to a Cepheid’s pulsations. When a star contracts, 
the gas beneath its surface gets hotter, but the extra heat does not raise the 
temperature; instead, it ionizes the helium. This ionized helium is very good at 
absorbing radiation. In other words, it becomes more opaque and absorbs the 
radiant energy flowing outward through it, toward the surface. This trapped heat 


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^ Gravity 


1 . Pressure forces exceed gravity and the pulsating 
star starts to expand 


2. Gravity and pressure balance each other, but 
inertia makes the star continue to expand 




Gravity 

3. Gravity exceeds the pressure 
and the star begins to contract 


4. Pressure and gravity once again achieve a 
balance, but the inertia makes the star 
continue to contract 



Figure 3.1 8 . Gravity and pressure during the pulsation cycle of a pulsating star. 


makes the star expand. This, then, provides the “push” that propels the surface 
layers of the star back outward. As the star expands, electrons and helium ions 
recombine, and this causes the gas to become more transparent (i.e., its opacity 
falls, and so the stored energy escapes). 

For a star to be susceptible to pulsations, it must have a layer beneath its 
surface in which the helium is partially ionized. The existence of such a layer 
will depend not only on the size and mass of a star, but also on its surface 
temperature, which, in most cases, will be in the range of 5000 to 8000 K. There 
is a region on the H-R diagram where such an area exists, and it is the location 
of the pulsating stars. It is called the instability strip. In this region are found 
the Cepheid variable and RR Lyrae stars. 



Stars 



3.13.2 Cepheid Variables and the 

Period-Luminosity Relationship 


Cepheid variables are named after 8 Cephei, which was the first star of its type to 
be discovered. It is a yellow giant star that varies by a factor of two in brightness 
over 5.5 days. 35 Figure 3.19 shows the variations of 8 Cephei in luminosity, size, 
and temperature. 

You will notice immediately that its luminosity and temperature have a 
maximum value when its size has a minimum value, and vice-versa; its size is at 
maximum when its luminosity and temperature are at minimum. Cepheids are 
very important for astronomers for two reasons. They can be seen at extreme 
distances, perhaps as great as a few million pc. This is because they are very 
luminous, with a range from a few hundred to a few tens of thousands of solar 
luminosity (i.e., 100 L Q to 10, 000 L Q ). Second, there exists a relationship between 
the period of a Cepheid and its average luminosity. The very faintest Cepheids 
(which are in fact hundreds of times brighter than the Sun) pulsate with a very 
rapid period of only one or two days, while the brightest (as much as 30,000 
times brighter than the Sun) have a much slower period of about 100 days. The 
correlation between the pulsation period and luminosity is called the period- 
luminosity relationship. If a star can be identified as a Cepheid, and its period 
measured, then its luminosity and absolute magnitude can be determined. This 
can then be used, along with its apparent magnitude, to determine its distance. 

The amount of metals in a Cepheid star’s outer layers will determine how 
it pulsates. This occurs because the metals can have a substantial effect on the 
opacity of the gas. They can then be classified according to their metal content. If 


Lui 



0 


7000 


Temperature ( K ) 


6000 



5000 



0 2 4 6 8 10 

Time (days) 


Figure 3.19. The size, temperature, and luminosity of 8 Cephei during one period. 



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Figure 3.20. Period-luminosity relationship for the two types of Cepheid variable star. 


a Cepheid is a metal-rich Population I star, 36 it is called a Type I Cepheid, and if it 
is a metal-poor Population II star, it is called a Type II Cepheid. Figure 3.20 shows 
a period-luminosity diagram for the two types of Cepheids. So, an astronomer 
must first determine what type of Cepheid he or she is observing before the 
period-luminosity relationship can be applied. 


3. 1 3.3 Cepheids: Temperature and Mass 

The period-luminosity relationship comes about because the more massive 
stars are also the most luminous stars as they cross the H-R diagram 
during core helium-burning. These massive stars are also larger in size and 
lower in density during this period of core-helium-burning, and the period 
with which a star pulsates is larger for lower densities; so, the massive 
pulsating stars have greater luminosities and longer periods. This is shown in 
Figure 3.21. 

We have seen that old, high-mass stars have evolutionary tracks that cross 
back and forth in the H-R diagram, and thus will intercept the upper end 
of the instability strip. Such stars become Cepheids when the helium ionizes 
at just the right depth to drive the pulsations. Those stars on the left (high 
temperature) of the instability strip will have helium ionization occurring too 
close to the surface and will involve only a small fraction of the star’s mass. 
The stars on the right (low-temperature) side will have convection in the star’s 
outer layers, and this will prevent the storage of the heat necessary to drive 
the pulsations. Thus, Cepheid variable stars can only exist in a very narrow 
temperature range. 



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Figure 3.21 . Instability strip and evolutionary tracks for stars of different mass. 


3.13.4 RR Lyrae and Long-Period 
Variable Stars 

The faintest and hottest stars on the instability strip are the RR Lyrae stars. These 
stars have a much lower mass than Cepheids. After the helium flash occurs, their 
evolutionary tracks pass across the lower end of the instability track as they move 
across the horizontal branch of the H-R diagram. The RR Lyraes, named after 
the prototype in the constellation Lyra, all have periods shorter than Cepheids, 
of about 1.5 hours to 1 day. They are small and dense stars compared with the 
Cepheids (but are nearly 10 times larger and about 100 times more luminous 
than the Sun). The RR Lyrae region of the instability strip is in fact a segment of 
the horizontal branch. They are all metal-poor Population II stars, and so many 
are found in globular clusters. 

The long-period variables are cool red giant stars that can vary by as much 
as a factor of 100, in a period of months or even years. Many have surface 
temperatures of about 3500 K and average luminosities in a range of 10 to as much 
as 10, 000L-,. They’re placed on the middle right-hand side of the H-R diagram. 
Many are periodic, but there are also a few that are not. A famous example of 
a periodic long-period variable star is Mira ( o Ceti) in Cetus. A famous non- 
periodic long-period variable star is Betelgeuse (a Orionis) in Orion. It may come 



1 20 Astrophysics is Easy 

as a surprise to you to know that we do not fully understand why some cool red 
giant stars become long-period variable stars. 

There are many pulsating stars that can be observed by the amateur 
astronomer, and in fact several organizations exist throughout the world that 
cater specifically to this pastime . 37 However, I shall just describe the brightest 
members of each of the three aforementioned classes: Cepheid, RR Lyrae, and 
long-period variable stars. All that is needed in observing these stars is a degree 
of patience, as the changes in magnitude can take as little as a few days to several 
hundred days, and, of course, clear skies. 

The nomenclature used in this list is the same as that used before, with a few 
changes. The apparent magnitude range of the variable is given, along with its 
period in days. 


3. 1 3.5 Bright Cepheid Variables 

8 Cephei HD 213306 22'’29.1 m +58°25' Jul-Aug-Sep 

3.48- 4.37 m — 3.32M 5.37 days F3-G3 Cepheus 

This is the prototype star of the classic short-period pulsating variables known 
as Cepheids. It was discovered in 1784 by the British amateur astronomer John 
Goodricke. It is a favorite with amateurs, as several bright stars also lie in 
the vicinity — Epsilon ( e ) Persei (4.2m), Zeta (£) Persei (3.4m), Zeta (£) Cephei 
(3.35m), and Eta ( 17 ) Cephei (3.43m). The behavior of the star is as follows: it 
will brighten for about 11/2 days and will then fade for 4 days, with a period 
of 5 days, 8 hours, and 48.2 minutes. Delta Cephei is also a famous double star, 
with the secondary star (6.3m) an attractive white color, which contrasts nicely 
with the yellowish tint of the primary. 

■») Aquilae HD 187929 19 ,! 52.5 m +01°00' Jun-Jul-Aug 

3. 48- 4. 39m -3.91M 7.17 days F6-G4 Aquila 

This is a nice Cepheid to observe, as its variability can be seen with the naked 
eye. The rise to brightest magnitude takes 2 days, and thereafter slowly fades. 
The nearby star Beta Aquilae (3.71 m) is often used as a comparison. It is the 
third-brightest Cepheid (in apparent magnitude), after Delta Cephei and Polaris. 
The actual period is 7 days, 4 hours, 14 minutes, and 23 seconds! 

RT Aurigae HD 45412 06 h 28.6 m +30°29' Nov-Dec-Jan 

5. 29-6. 6m -2.65M 3.73 days F5-G0 Auriga 

Also known as 48 Aurigae, this star was discovered to be a variable in 1905 by T. 
Astbury, who was a member of the British Astronomical Association. The rise to 
maximum takes IV 2 days, with a diminishing over 2Vi days. Easy to observe in 
binoculars, it lies midway between Epsilon (s) Geminorum (3.06 m) and Theta 
( 6 ) Auragae (2.65 m). The period has been measured to an astounding accuracy 
to be 3.728261 days! 

a Ursae Minoris ADS 1477 02 , ’31.8 m +89°16' Sep— Oct— Nov 

1.92-2. 07m — 3.64M 3.97 days F7:Iib-Iiv Ursa Minor 


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121 


Possibly the most famous star in the sky. Polaris, or the Pole star, is located 
less than a degree from the celestial pole. It is one of the Type II Cepheids, 
which are also known as W Virginis variable stars. The magnitude changes are 
very small and therefore not really detectable with the naked eye. Polaris also 
consists of double color, namely a yellowish primary and a faint whitish-blue 
secondary at a magnitude of 8.2. The primary is also a Population II Cepheid 
variable and a spectroscopic binary. Although claims have been made to the 
effect that the system can be resolved in an aperture as small as 4.0 cm, at least 
6.0 cm will be required to split it clearly. Will be closest to the actual pole in 
2102 A.D. 

Other Cepheid variable stars that can be observed with amateur equipment 
are: U Aquilae, Y Ophiuchi, W Sagittae, SU Cassiopeiae, T Monocerotis, and T 
Vulpeculae. 

3. 1 3.6 Bright RR Lyrae Variables 

RR Lyrae HD 182989 19 h 25.3 m +42°47' Jun-Jul-Aug 

7. 06-8. 12m 1.13 M 0.567 days see text below* Lyra 

This is the prototype of the RR Lyrae class of pulsating variable stars. These are 
similar to Cepheids but have shorter periods and lower luminosities. There are 
no naked-eye members of this class of variable, and RR Lyrae is the brightest 
member. There is a very rapid rise to maximum, with the light of the star 
doubling in less than 30 minutes, with a slower falling in magnitude. From an 
observational viewpoint, it is a nice white star, although detailed measurements 
have shown that it does become blue as it increases in brightness. There is 
some considerable debate as to the changes in spectral type that accompany the 
variability. *One source quotes A8-F7, while the other is A2-F1. Take your pick. 
There is also some indication that there is another variability period along with 
the shorter one, which lasts about 41 days. 

Other RR Lyrae variable stars are: RV Arietis, RW Arietis, and V467 Sagittari; 
however, all of these stars are faint and so will present a considerable challenge 
to observers. 


3.13.7 Long-Period Variables 

Mira o Cet 02 , '19.3 m -02°59' Sep-Oct-Nov 

2.00-10m — 3.54M 331.96 days M9-M6e Cetus 

An important star, and maybe the first variable star ever observed. Written 
records certainly exist as far back as 1596. The prototype of the long-period 
pulsating variable, it varies from 3rd to 10th magnitude over a period of 332 days, 
and it is an ideal star for the first-time variable star observer. At minimum, the 
star is a deeper red color, but, of course, fainter. It now has a lower temperature 
of 1900 K. The period, however, is subject to irregularities, as is its magnitude, 
and can be longer, or shorter, than the quoted average of 332 days. It has been 
observed for maximum light to reach 1st magnitude, similar to Aldebaranl One 


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of the oddities about Mira is that the change in spectral class does not occur 
exactly with maximum, but rather a few days later! Another oddity is that when 
Mira is at its faintest, it apparently is also at its largest, when you would think 
the opposite to be true. A reason for this has been put forward recently: the 
star produces titanium oxide in its atmosphere as it cools and expands. The 
compound then acts as a filter, blocking out the light. The name Mira is Arabic 
for “wonderful star.” 

Other Mira-type stars are R Leonis and R Leporis, both of which have been 
described in earlier sections. 


3.14 The Death of Stars 


Stars live for millions, billions, and even hundreds of billions of years, 38 and 
so you may be thinking how on Earth can we know anything about how a star 
dies? After all, we have only been on a planet that is about 4.5 billion years old 
and studying astronomy for about 10,000 years. Well, fortunately for us, it is 
nevertheless possible to observe the many disparate ways in which a star can end 
its life. 

Once again, it is the mass of a star that decides how it will end its life, 
and the results are spectacular and sometimes very strange indeed. Low-mass 
stars can end their lives in a comparatively gentle manner, forming beautiful, 
and apparently delicate structures that we know as planetary nebulae, before 
proceeding to small and ever-cooling white dwarf stars. At the other end of the 
scale, high-mass stars tend to end their lives in a far more spectacular fashion 
by exploding! These are the rare supernovae. 

We begin our journey by looking at stars that have a low mass. 


3. 1 5 The Asymptotic Giant Branch 


Let’s recap briefly how low-mass stars (and by this I mean stars with a mass 
of about 4 M q and less), behave after leaving the main sequence. When core 
hydrogen-burning ceases, the core will shrink, and this heats up the surrounding 
hydrogen gas, and so hydrogen-shell burning begins. The outer layers of the 
star will expand but also cool, and so the star becomes a red giant. The post- 
main-sequence star will move up and to the right on the H-R diagram as its 
luminosity increases and temperature falls. We can say that the star now lies 
on the red-giant branch of the H-R diagram. The next stage involves the onset 
of helium-burning in the core. If a star has a high mass (greater than about 
2-3 M q ), then this starts gradually, but if the star has a lower mass, this stage 
begins suddenly, in what is called the helium flash. But no matter which way 
it starts, a result of the helium-burning is that the core actually cools down, 
with a resulting slight decrease in luminosity. The outer layers of the star also 
contract a little, heating up in the process, and so the evolutionary track of 
the red giant now moves left across the H-R diagram. The luminosity during 





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this phase remains more or less constant, so the path is nearly horizontal, and 
so is called the horizontal branch. Stars on the horizontal branch are stars in 
which helium-burning is occurring in the core, which in turn is surrounded 
by a shell of hydrogen-burning. Many such stars are often found in globular 
clusters. 

We can now look at the next stage of a star’s life. Recall from Section 3.10.1 
that by-products of the triple a process are the elements carbon and oxygen. So 
after a suitably long period of time, maybe 100 million years, we could expect 
all of the helium in the core to have been converted into carbon and oxygen. 
This would mean that core-helium-burning would cease. A similar process to 
that (which was explained in Section 3.10.2) then begins. The absence of nuclear 
fusion results in a contraction of the core because there is no energy source 
to provide the internal pressure necessary to balance the force of gravity. The 
core contraction is stopped, however, by degenerate electron pressure, which 
is something we met earlier. A result of the core-contraction is a release of 
heat into the helium gas surrounding the core, and so helium-burning begins 
in a thin shell around the carbon-oxygen core. This is aptly called shell helium- 
burning. 

Now an extraordinary thing happens: the star enters a second red-giant phase. 
It is as if history has repeated itself. Stars become red giants at the end of their 
main-sequence lifetimes. The shell hydrogen-burning phase provides energy, 
causing the outer layers of the star to expand and cool. In a similar fashion, the 
energy from the helium-burning shell also causes the outer layers to expand, and 
so the low-mass star rises into the red-giant region of the H-R diagram for a 
second time. But this time it has an even greater luminosity. 

This phase of a star’s life is often called the asymptotic giant branch phase, 
or AGB. Thus, these stars are called AGB stars and are on the asymptotic giant 
branch of the main sequence. 

The structure of an AGB star is shown in Figure 3.22. Its central region is 
a degenerate carbon-oxygen mix surrounded by a helium-burning shell, which 
in turn is surrounded by a helium-rich shell. This is further surrounded by a 
hydrogen-burning shell. All of this is further encompassed by a hydrogen-rich 
outer layer. What is truly remarkable is the size of these objects. The core region 
is about the same size as the Earth, while the hydrogen envelope is immense. 
It can be as large as the orbit of the Earth! When the star has aged, however, 
the outer layers, which are expanding, cause the hydrogen-burning shell to also 
expand and thus cool, and so the nuclear reactions occurring therein may cease, 
albeit temporarily. 

The luminosity of these stars can be very high indeed. For example, a 1 M Q 
AGB star may eventually attain a luminosity of 10, 000 L 0 . Compare this with 
the luminosity of only 1000 L Q it is achieves when it reaches the helium-flash 
phase, and the poor 1 L 0 when it resides on the main sequence. It is sobering to 
think ahead and imagine what will happen when the Sun becomes an AGB star 
in about 8 billion years from now! 

There are many AGB stars that can be observed by amateurs, and in fact 
we have already mentioned and described several of them: the archetypal AGB 
star is Mira (o Ceti), but there are also R Leonis, R Leporis, R Aquarii, and 


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Hydrogen-burning shell 



Figure 3.22. The structure of an AGB star. 


R Cassiopea. In addition, there are a few others, although they are somewhat 
fainter. These include \ Cygni, W Hydrae, S Pegasi, and TT Monocerotis. 


3.16 Dredge-Ups 


We have seen that energy and heat are transported from a star’s core to the 
surface by two methods: convection and radiation. Convection is the motion of 
the star’s gases moving upwards toward the surface, and then cooling of the gases 
so that it falls downwards. This method of energy transfer is very important 
in giant stars. Radiation, or radiative diffusion as it is also sometimes called, 
is the transfer of energy using electromagnetic radiation, and is only important 
when the gases in a star are transparent (the opacity is low). When a star ages 
and leaves the main sequence, the convective zone can increase substantially in 
size, and sometimes extend right down to the core. This means that the heavy 
elements, or metals, that are formed there can be carried to the star’s surface by 
convection. This process has the very unglamorous name of dredge-up. Th e first 
dredge-up begins when the star becomes a red giant for the first time (i.e., core 
hydrogen-burning phase has stopped). The by-products of the CNO 39 cycle of 
hydrogen are transported to the surface because the convective zone now reaches 
deep into the core regions. A second dredge-up begins when the helium-burning 
phase ends. Then, during the AGB phase, a third dredge-up occurs, but only if 
the mass of the star is greater than 2M 0 , when a large amount of newly formed 
carbon is carried to the star’s surface. The spectrum of a star that has such a 
carbon-enriched surface exhibits very prominent absorption bands of carbon- 
rich elements, such as C 2 , CH, and CN. Such stars that have undergone a third 
dredge-up are often called carbon stars. 



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125 


3.1 7 Mass Loss and Stellar Winds 


As a star continues to rise up the AGB, it increases in both brightness and size, and, 
consequently, it develops a very strong stellar wind. This blows the star’s outer layers 
into interstellar space. Thus, the star undergoes a substantial mass loss during this 
phase, maybe as much as 10 -4 M Q per year (this is about 1000 times greater than the 
mass loss of a red giant star, and about 10 billion times the mass loss of the present- 
day Sun). The cause of these extreme stellar winds is still a puzzle, although the 
surface gravity of AGB stars is very lowbecause the stars are so large; thus, any sort of 
disturbance on the star surface is capable of expelling material outwards. The outer 
layers of the star flow outward at 10 km per second (about 2% of the speed of the 
solar wind), cooling as they move from the star. Dust particles can thus form in the 
cooler surrounding gas formed out of the ejected carbon-rich molecules. In fact, it is 
believed that tiny grains of soot are formed! Many carbon stars have been observed, 
surrounded by cocoons of carbon-rich matter. In some cases, the dust cloud is so 
thick that it can totally obscure the star, absorbing all the emitted radiation. The 
dust then heats up, and re-emits the energy, but this time in the infrared. 


3.18 Infrared Stars 


It may come as a surprise to know that AGB stars, which can have luminosities 
10,000 times that of the Sun, were, until the 1960s, hardly known. The reason 
for this is simple: the dust that surrounds the star, and re-emits the radiation, 
is so cool that the reradiated energy is almost entirely in the infrared part of 
the spectrum. This is, of course, invisible to the naked eye, and it has only been 
explored in detail in the past 30 years. This is shown in Figure 3.23. These stars 


Brightness 



Figure 3.23. The spectra of an infrared star compared with the Sun. 




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are very faint, or even invisible in the visible part of the spectrum. By comparison, 
the Sun is very bright in the visible part of the spectrum, and very faint in the 
infrared. The surface of an infrared star can be thought of as starting at the 
surface layer of the dust cloud, and for some AGB stars, this can have a radius 
as much as 500 AU, which is about 10 times the size of the Solar System. These 
outer layers of the star are extremely tenuous and hold only a fraction of its 
total mass. The vast majority of the mass is in the carbon-oxygen core and the 
energy-forming layers that surround it. So, we can picture an infrared star as 
having a very small and dense central part and an enormous, low-density outer 
layer. 

Most of the energy emitted by the Sun is in the visible part of the spectrum. 
In comparison, nearly all of the energy radiated from the dust surrounding an 
infrared star is invisible to the naked eye. 


3.19 The End of an AGB Star's Life 


As it ages, an AGB star continues to grow in size and increase its luminosity, 
along with an increase in the rate at which it loses mass. As mentioned earlier, 
the mass loss can be 10 -4 M o per year, which means that if the Sun lost mass 
at this rate, it would only last for 10,000 years. So, obviously, even giant stars 
cannot carry on this way for very long. If a star has a mass of less than about 
8M 0 , its stellar wind will soon strip away the outer layers almost down to the 
degenerate core. Therefore, a loss of the outer layers would signal the end of the 
AGB phase. For stars that are greater than 8M 0 , the end of the AGB phase arrives 
in a much more spectacular event — a supernova... which will be discussed in a 
later section. 

I would like to end this section on the AGB part of a star’s life with a sobering 
and amazing thought. Carbon stars enrich the interstellar medium with not only 
carbon, but with some nitrogen and oxygen, as well. In fact, carbon can only be 
formed by the triple a process that occurs in helium-burning, and carbon stars 
are the main means by which the element carbon is dispersed throughout the 
interstellar medium. The part that always amazes me is when I consider that the 
carbon in my body, and in fact in the body of every living creature on the Earth, 
was formed many billions of years ago, inside a giant star undergoing the triple 
a process. It was then dredged up to the star’s surface and expelled into space. 
Later, by some means, it formed the precursor to the Solar System and made the 
Sun and planets, and all life on the Earth. 

We are made of the stuff of stars! 

One of the benefits of a carbon star, from the point of view of an observer, 
is that many of the rare carbon stars are visible in the night sky with amateur 
instruments. We have already come across a few of these: R Leporis, RS Cygni, 
and 19 Piscium. But there are several more that are worth seeking out, and I have 
listed those on the following pages. The one aspect of these stars that will be 
immediately apparent to you is their color; all of them are strongly red colored. 
They are, in fact, the reddest stars visible to the amateur astronomer. 


Stars 


3. 1 9. 1 Bright Carbon Stars 

XCnc HD 76221 08 h 55A m +17°14' Jan-Feb-Mar 

6.12 v m B— V :2.97 C6 Cancer 

An extremely orange star, this semi-regular variable star, classification SRB, has 
a period of 180 to 195 days and has been observed to range in magnitude from 
5.6 to 7.5. 

La Superba Y CVn 12 ft 45.1 m +45°26' Mar-Apr-May 

5.4 v m B— V:2.9 C7 Canes Venatici 

The color of this star (red) is best seen through binoculars or a small telescope. 
With a period of 159 days and varying in magnitude between 4.9 and 6.0 m, this 
red giant has a diameter of 400 million kilometres. 

V Pav HD 160435 17 h 43.3 m -57° 43' May-Jun-Jul 

6.65 v m B— V:2.45 C5 Pavo 

A red giant variable star, class SRB, varying in brightness from 6.3 to 8.2 m over 
a period of 225.4 days. It also has a secondary period of about 3735 days. A 
glorious deep-red color. 

V Aql HD 177336 19''04.4 m -05°41' Jun-Jul-Aug 

7.5 v m B— V:5.46 C5 Aquila 

A semi-regular variable star with a period of about 350 days, varying in magnitude 
from 6.6 to 8.1 m. A very deep red in color. 

S Cephei HD 206362 2l'’35.2 m +78°37' Jul-Aug-Sep 

7.9 v m B— V:2.7 C6 Cepheus 

A moderately difficult star to observe due to its magnitude range of 7 to 12 m; 
nevertheless, it has a very high color index, making it one of the reddest stars in 
the sky, if not the reddest. Its red color immediately strikes you and, once seen, 
is never forgotten. 

R Scl HD 8879 01 , '26.9 m -32°33' Sep-Oct-Nov 

5.79 v m B— V:1.4 C6 Sculptor 

A semi-regular-period variable star, with a period ranging between 140 and 146 
days; it varies in brightness from 5.0 to 6.5. 

U Cam 03*41. 8 m +62°39' Oct-Nov-Dec 

8.3 v m B— V:4.9 N7 Camelopardalis 

A semi-regular variable star, period 412 days with a magnitude range of 7.7 to 
9.5 m. It has a very deep red color. 

W Ori HD 32736 05'’0.4 m +01° 11' Nov-Dec-Jan 

6.3 v m B— V:3.33 N5 Orion 


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A red giant variable star, classification SRB, with a period of 212 days, although 
a secondary period of 2450 days is believed to occur. Varies in magnitude from 
5.5 to 7.7m. A deep-red star. 

R Corona Borealis HD 141527 15 h 48.6 m +28°09' Apr-May-Jun 

5.89 v m B— V:0.608 GOI ab:pe Corona Borealis 

Although not strictly a carbon star, R Cor Bor, as it is known affectionately, should 
be mentioned here. It is the prototype variable star of the class RCB. What makes 
this star so special is that it is an irregular variable, usually seen at maximum 
brightness, but then suddenly fading down to 12th magnitude, which can last 
for several weeks, months, or even as long as a year. 40 Then, just as suddenly, 
it can return to its normal brightness. The reason for this strange behavior is 
that carbon grains condense out in the star’s atmosphere, thus blocking out the 
light from the star. Radiation then causes the grains to dissipate, and so the 
star returns to its usual magnitude. The cycle then begins again with the grains 
building up over time. Other stars that show a similar behavior are RY Sagittarii 
(6.5 m), SU Tauri (10 m), and S Apodis (10 m). 


3.20 Planetary Nebulae 


At the end of the AGB phase, all that will remain of a star is the degenerate core 
of carbon and oxygen, surrounded by a thin shell in which hydrogen-burning 
occurs. The dust ejected during the AGB phase will be moving outward at tens of 
kilometers per second. As the debris moves away, the hot, dense, and small core 
of the star will become visible. The aging star will also undergo a series of bursts 
in luminosity, and during each burst, eject a shell of material into interstellar 
space. The star now begins to move rapidly toward the left of the H-R diagram, 
at an approximately constant luminosity but at an increasing temperature. It 
will only take, say, a few thousand years for the surface temperature to reach 
30,000 K. Some stars achieve temperatures of 100,000 K. At these high tempera- 
tures, the exposed core of the star will emit prodigious amounts of ultraviolet 
radiation, which can excite and ionize the expanding shell of gas. The shell of 
ionized and heated gas will begin to glow and produce what is called a planetary 
nebula. 

We know that as the helium in the helium-burning shell is depleted, the 
pressure that supports the dormant hydrogen-burning shell decreases. Therefore, 
the hydrogen-burning shell contracts and heats up, thereby initiating hydrogen- 
burning. This newly started hydrogen-burning creates helium, which falls down 
upon the temporarily dormant helium-burning shell. If the shell temperature 
reaches a specific value, it reignites in what is called the helium-shell flash, 
similar to (but less intense than) the helium flash that occurs in the evolution 
of low-mass stars. The newly created energy pushes the hydrogen-burning shell 
outward, cooling it as it does so, which results in a cessation of the hydrogen- 
burning, and the shell becomes dormant once again. The process then starts all 
over again. 



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129 


The luminosity of the AGB star increases quite substantially when the helium 
shell flash occurs, although it is only for a relatively short time. This short- 
lived burst is called a thermal pulse. After a thermal pulse has occurred, the 
star resumes its former appearance until enough helium builds to allow another 
thermal pulse to occur. With each thermal pulse, the mass of the degenerate core, 
consisting of carbon and oxygen, will increase. For the very massive stars, the 
thermal pulse occurs in the very deep interior of the star and produces only a 
slight, temporary change in luminosity. For a star of mass 1 M 0 , a thermal pulse 
would be close enough to the surface to cause the luminosity to increase by a 
factor of 10 and last about 100 years. The time between thermal pulses varies 
depending on the star’s mass, but calculations predict that they would occur at 
ever-decreasing intervals, perhaps as short as 100,000 to 300,000 years, while the 
luminosity of the star during this time would slowly increase overall. 

Significant mass loss can also occur during thermal pulses. A star’s outer 
layers can separate completely from the carbon- and oxygen-rich core, and as 
the ejected material disperses into space, grains of dust can condense out of the 
cooling gas. The radiation from the very hot core can propel the dust grains 
farther, and so the star sheds its outer layers completely. In this manner, a star 
of mass, say, 1 M 0 can lose about 40% of its mass. Even more mass is lost by 
the massive stars. As the dying star loses its outer layers, the hot core is exposed 
and it illuminates the surrounding dust and gas cloud. 

The evolution of the remaining core is itself of interest as it progresses rapidly 
to its final state. There are two factors that can influence the rate at which the 
core evolves. First, due to the star’s extreme luminosity (which can be as high as 
100, 000 L Q ), it consumes its hydrogen at a very fast rate. Second, little hydrogen 
remains in the thin hydrogen-burning shell that surrounds the degenerate core, 
so there is hardly any fuel left to be consumed. The central stars of some planetary 
nebulae have as little as a few millionths of a solar mass of hydrogen left to 
burn, and so they fade very rapidly. In fact, some can have their luminosities 
decrease by as much as 90% in as little as 100 years, whereas others may require 
a bit longer, perhaps a few thousand years. As the source of ionizing photons 
decreases over time, the planetary nebulae grow darker and eventually fade 
away. 

Planetary Nebulae 41 are some of the most interesting and beautiful objects 
in the sky, and they have a lot to offer to the amateur. They range across the 
whole of the observational spectrum: some are easy to find in binoculars, while 
others require a large aperture, patience, and maybe even specialized filters to 
be distinguished from the background star fields. These small shells of gas, the 
atmosphere of stars come in a variety of shapes, sizes, and brightnesses. Many 
have a hot central star within the nebula, which is visible in amateur equipment 
and is the power source, providing the energy for the gas to glow. 

Several nebulae have a multiple-shell appearance, and this is thought to be 
due to the red giant experiencing several periods of pulsation where the material 
escapes from the star. The strong stellar winds and magnetic fields of the star 
are also thought to be responsible for the many observed exotic shapes of the 
nebulae. Planetary nebulae are only a fleeting feature in our Galaxy; after only a 
few tens of thousands of years, they will have dissipated into interstellar space, 
and so no longer exist. Thus, the planetary nebulae we observe today cannot 


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be older than about 60,000 years. However, this aspect of a star’s evolution is 
apparently very common, and there are more than 1400 planetary nebulae in our 
part of the Galaxy alone! 

Visually, the nebulae are one of the few deep sky objects that actually appear 
colored. About 90% of their light comes from the doubly-ionized oxygen line, 
OIII, at wavelengths 495.9 nm and 5000.7 nm. This is characteristic of blue-green 
color and, it so happens, the color at which the dark-adapted eye is at its 
most sensitive. Specialized light filters are also extremely useful for observing 
planetaries, as they isolate the OIII light in particular, increasing the contrast 
between the nebulae and the sky background, thus markedly improving the 
nebulae’s visibility. 

Such is the variety of shapes and sizes that there is something to offer all 
types of observers. Some planetaries are so tiny that even at high magnification, 
using large-aperture telescopes, the nebulae will still appear starlike. Others are 
much larger. For instance, the Helix Nebula, Caldwell 63, is half the size of the 
full Moon but can only be observed with low magnification and perhaps only in 
binoculars, as any higher magnification will lower its contrast to such an extent 
that it will simply disappear from view. Many, such as the Dumbbell Nebula, 
M27, in Vulpecula, exhibit a bipolar shape. Still others show ring shapes, such as 
the ever-popular Ring Nebula, M57, in Lyra. 

An interesting aspect is the possibility of observing the central stars of the 
nebulae. These are very small subdwarf and dwarf stars. They are similar to 
main-sequence stars of types O and B, but, as they are running down their 
nuclear reactions, or in some cases, no longer producing energy by nuclear 
reactions, they are consequently fainter and smaller. These two characteristics 
make observation very difficult. The brightest central star is possibly that of 
NGC 1514 in Taurus, at 9.4 magnitude, but the majority are at magnitude 10 or 
fainter. 

There is a classification system called th eVorontsoz-Vellyaminov Classification 
System, which can be used to describe the appearance of a planetary. Although 
it is of limited use, it will be used here. 


Planetary Nebulae Morphology Types 


1 Starlike 

2 Smooth disc-like appearance 

a. bright toward center 

b. uniform brightness 

c. possible faint ring structure 

3 Irregular disc-like appearance 

a. irregular brightness distribution 

b. possible faint ring structure 

4 Definite ring structure 

5 Irregular shape 

6 Unclassified shape* 

* can be a combination of two classifications, 
(e.g., 4 + 3, ring and irregular disc) 


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131 


The usual information is given for each object, with the addition of morphology 
class [O] and central star brightness [#]. In addition, the magnitude quoted is 
the magnitude of the planetary nebula as if it were a source point. This last 
parameter can often be confusing, so even if a nebula has a quoted magnitude 
of, say, 8, it may be much fainter than this and, consequently, hard to find. 

3.20.1 Bright Planetary Nebulae 

Caldwell 39 NGC 2392 07 h 29.2 m +20°55' Dec-Jan-Feb 
8.6m ©15" ©3b+3b *119.8 Gemini 

Also known as the Eskimo Nebula. This is a small but famous planetary nebula, 
which can be seen as a pale blue dot in a telescope of 10 cm, although it can be 
glimpsed in binoculars as the apparent southern half of a double star. Higher magni- 
fication will resolve the central star and the beginning of its characteristic “Eskimo” 
face. With an aperture of 20 cm, the blue disc becomes apparent. Research indicates 
that we are seeing the planetary nebula pole-on, although this is by no means 
certain. Its distance is also uncertain, with values ranging from 1600 to 7500 l.y. 

Caldwell 59 NGC 3242 10 h 24.8 m -18°38' Jan-Feb-Mar 
8.4m ©16" ©4+3b #12.1 Hydra 

Also known as the Ghost of Jupiter. One of the brighter planetary nebulae and 
the brightest in the spring sky for northern observers, this is a fine sight in 
small telescopes. Visible in binoculars as a tiny blue disc. With an aperture of 
10 cm, the blue color becomes more pronounced along with its disc, which is 
approximately the same size as that of Jupiter in a similar aperture. The central 
star has a reported temperature of about 100,000 K. 

Messier 97 NGC 3587 ll h 14.8 m +55°01' Feb-Mar-Apr 
9.9m © 194" ©3a #16 Ursa Major 

Also known as the Owl Nebula. Not visible in binoculars due to its low surface 
brightness; apertures of at least 20 cm will be needed to glimpse the “eyes” of the 
nebula. At about 10 cm aperture, the planetary nebula will appear as a very pale 
blue-tinted circular disc, although the topic of color in regard to this particular 
planetary nebula is in question. 

Caldwell 6 NGC 6543 17 h 58.6 m +66°38' May-Jun-Jul 

8.3m ©18|350" ©3a+2 #11 Draco 

Also known as the Cat’s Eye Nebula. This is seen as a bright oval planetary 
nebula with a fine blue-green color. Caldwell 6 is one of the planetary nebula 
that became famous after the HST published its image. Visible even in a telescope 
of 10 cm, but a large telescope (20 cm) will show some faint structure, while to 
observe the central star requires a 40 cm aperture. The incredibly beautiful and 
complex structure is thought to be the result of a binary system, with the central 
star classified as a Wolf-Rayet star. 

Messier 57 NGC 6720 18 h 53.6 m +33°02' Jun-Jul-Aug 

8.8m ©71' 04+3 #15.3 Lyra 


Astrophysics is Easy 


132 

Also known as the Ring Nebula. The most famous of all planetary nebulae, 
surprisingly — and pleasantly — visible in binoculars. However, it will not be 
resolved into the famous “smoke-ring” shape seen so often in color photographs; 
it will, rather, resemble an out-of-focus star. It is just resolved in telescopes 
of about 10 cm aperture, and at 20 cm the classic smoke-ring shape becomes 
apparent. At high magnification (and larger aperture), the Ring Nebula is truly 
spectacular. The inner region will be seen to be faintly hazy, but a large aperture 
and perfect conditions will be needed to see the central star. 

Caldwell 15 NGC 6826 19 h 44.8 m +50°31' Jun-Jul-Aug 

8.8m ©25" ©3a+2 #11 Cygnus 

Also known as the Blinking Planetary. A difficult planetary nebula to locate, 
but well worth the effort. The blinking effect is solely due to the physiological 
structure of the eye. If you stare at the central star long enough, the planetary 
nebula will fade from view. At this point, should you move the eye away from 
the star, the planetary nebula will “blink” back into view at the periphery of your 
vision. Although not visible in amateur telescopes, the planetary nebula is made 
up of two components: an inner region consisting of a bright shell and two ansae 
(delicate protuberances from either side), and a halo that is delicate in structure 
with a bright shell. 

Messier 27 NGC 6853 19 h 59.6 m +22°43' Jun-Jul-Aug 

7.3m © 348" 0 3+2 #13.8 Vulpecula 

Also known as the Dumbbell Nebula. This famous planetary nebula can be seen 
in small binoculars as a box-shaped hazy patch, and many amateurs consider it 
the sky’s premier planetary nebula. In apertures of 20 cm, the classic dumbbell 
shape is apparent, with the brighter parts appearing as wedge shapes that spread 
out to the north and south of the planetary nebula’s center, and a central star 
may be glimpsed. 

Herschel 16 NGC 6905 20 h 22.4 m +20°05' Jun-Jul-Aug 

11.1m ©40" 03+3 #15.5 Delphinus 

Also known as the Blue Flash Nebula. The true nature of this planetary nebula 
only becomes apparent at apertures of at least 20 cm, when the lovely blue color 
is seen. The central star can be seen only under good seeing conditions. 

Caldwell 55 NGC 7009 21 h 04.2 m -11°22' Jul-Aug-Sep 

8.3m © 25" 04 + 6 #12.78 Aquarius 

Also known as the Saturn Nebula. Although it can be glimpsed in small apertures, 
a telescope of at least 25 cm is needed to see the striking morphology of the 
planetary nebula that gives it its name. There are extensions, or ansae, on either 
side of the disc, along an east-west direction, which can be seen under perfect 
seeing conditions. High magnification is also justified in this case. Recent theory 
predicts a companion to the central star, which may be the cause of the peculiar 
shape. 

Caldwell 63 NGC 7293 22 h 29.6 m -20° 48' Jul-Aug-Sep 

6.3m © 770" 04+3 #13.5 Aquarius 


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133 


Also known as the Helix Nebula. Thought to be the closest planetary nebula to 
the Earth, at about 450 l.y., it has an angular size of over 1/4° — half that of the 
full Moon. However, it has a very low surface brightness, and is thus notoriously 
difficult to locate. With an aperture of 10 cm, low magnification is necessary, and 
averted vision is useful to glimpse the central star. The use of an OIII filter will 
drastically improve the image. 

Caldwell 22 NGC 7662 23 h 25.9 m +42°33' Aug-Sep-Oct 
8.6m © 12" 04+3 #13.2 Andromeda 

Also known as the Blue Snowball. This planetary nebula is visible in binoculars 
due to its striking blue color, but it will only appear stellar-like. Research indicates 
that the planetary nebula has a structure similar to that seen in the striking 
HST image of the Helix Nebula, showing Fast Low-Ionization Emission Regions 
(Fliers). These are clumps of above-average-density gas ejected from the central 
star before it formed the planetary nebula. 

Messier 76 NGC 650 01 h 42.4 m +51°34' Sep-Oct-Nov 
10.1m © 65" 0 3+6 #15.9 Perseus 

Also known as the Little Dumbbell Nebula. This is a small planetary nebula that 
shows a definite non-symmetrical shape. In small telescopes of aperture 10 cm, 
and using averted vision, two distinct “nodes” or protuberances can be seen. 
With apertures of around 30 cm, the planetary nebula will appear as two bright 
but small discs that are in contact. Even larger telescopes will show considerably 
more detail. 

Herschel 53 NGC 1501 04 h 07.0 m +60°55' Oct-Nov-Dec 
11.5m © 52' 0 3 #14.5 Camelopardalis 

Also called the Oyster Nebula. A blue planetary nebula, easily seen in telescopes of 
20 cm and glimpsed in apertures of 10 cm. With a larger aperture, some structure 
can be glimpsed, and many observers liken this planetary nebula to that of the 
Eskimo Nebula. 


3.21 White Dwarf Stars 


We now look at the endpoint for low-mass stars, and it is a very strange end 
indeed. We have seen that stars with a mass of less than 4M Q never manage to 
produce the internal pressure and temperature necessary to provide the means 
to burn the carbon and oxygen in the core. What happens instead is an ejection 
of the star’s outer layers, leaving behind the very hot carbon-oxygen-rich core. 
In such a scenario, the core has stopped producing energy by nuclear fusion and 
so just cools down, admittedly over a vast time scale. These cooling relics are 
called white dwarf stars. In many instances, they are no bigger than the Earth. 




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3.21.1 Electron Degeneracy 

Experience tells us that as the mass of an object increases, so does its size, and 
this applies for many astronomical objects, such as stars on the main sequence. 
However, the opposite is true for white dwarfs. The more massive a white dwarf 
star, the smaller it is. This contrary behavior has to do with the electron structure 
of the material of the white dwarf. Increasing the density of an object will lead to 
an increase in pressure, as observed in main-sequence stars, but the pressure in 
a white dwarf star (which is, remember, the core of a once-much-larger star) is 
produced by degenerative electrons. 42 This electron degenerate pressure supports 
the star. An increase in density, however, also leads to an increase in gravity. For 
the white dwarf star, this increased gravity will exceed the increase in pressure, 
and so the star will contract. As it gets smaller, both the gravity and pressure 
increase further and come into balance with each other, but at a smaller size for 
the white dwarf. This means the more massive a white dwarf star, the smaller it 
is. As an example, a 0.5 M 0 white dwarf star is about 90% larger than the Earth, 
whereas a 1 M Q white dwarf star is only about 50% larger than the Earth. If the 
white dwarf is 1.3 M Q , then it is only 40% as large as the Earth. 

3.21 .2 The Chandrasekhar Limit 

White dwarf stars have a very unusual mass-radius relationship, shown in 
Figure 3.24. As you can see, the more degenerate matter you put into a white 
dwarf star, the smaller it gets. However, you cannot do this ad infinitum, as 
there is a maximum mass that a white dwarf can have. This mass, which is about 
1.4M 0 , is called the Chandrasekhar limit, named after the Indian scientist who 
first seriously studied the behavior of white dwarfs. It is the mass for which the 



Figure 3.24. Mass-radius relationship for white dwarf stars. 



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135 


mass-radius relationship drops to zero, so that a white dwarf star with a mass 
equal to the Chandrasekhar limit will shrink to a very small size. But no star 
with a mass greater than 1.4M Q can be supported against the crush of gravity 
by the pressure of the degenerate electrons. This means that the main-sequence 
stars of types O, B, and A, which have masses greater than the Chandrasekhar 
limit, will need to shed mass if they are to become white dwarf stars. This they 
do while becoming AGB stars, as we saw earlier. But not all stars do achieve the 
necessary mass loss, and in such cases where the contraction cannot be stopped 
by degenerate electrons, the stars collapse even further to become neutron stars, 
and perhaps even black holes. 

The radius is given in terms of the Earth’s radius. The more massive a white 
dwarf, the smaller it will be. Note that on this graph, the size of a white dwarf 
will fall to zero if it has a mass of 1.4M Q . 

A question that is often asked is, “What is a white dwarf star made of?” The 
answer is surprising. The matter making up the white dwarf star consists mostly 
of ionized oxygen and carbon atoms, which are floating in a sea of fast-moving 
degenerate electrons. As the star continues to cool, the particles in this matter 
slow down, resulting in electric forces between the ions beginning to dominate 
the random thermal motions they may have originally had. These ions no longer 
move freely through the white dwarf but are aligned in orderly rows, rather 
like a giant crystal lattice. It is appropriate to think of the white dwarf as being 
“solid,” with the degenerate electrons still moving freely in the crystal lattice, 
just as electrons move in, say, a copper wire. Another interesting point to make 
is that a diamond is a crystal lattice of carbon, so a cooling white dwarf star can 
also be thought of as a (sort of) giant spherical diamond. The density in a white 
dwarf star is immense, typically 10 9 kg m . This is about one million times the 
density of water. One of the statistics astronomers like to throw out is that one 
teaspoon of white dwarf matter weighs about 5.5 tons, equal to the weight of an 
elephant. . .providing, of course, that you could get a teaspoon of the matter to 
the Earth in the first place. 


3.21.3 White Dwarf Evolution 

When a white dwarf shrinks to its ultimate size, it will no longer have fuel 
available for nuclear fusion. It will, however, still have a very hot core and a 
large reservoir of residual heat. For example, the surface temperature of a famous 
white dwarf star, Sirius B, is about 30,000 K. Time will pass, and with it, the 
white dwarf will cool down as it radiates its heat into space. As it does so, it will 
also grow dimmer, as shown in Figure 3.25, where white dwarf stars of differing 
mass are plotted on an H-R diagram. The more massive a white dwarf star, the 
smaller its surface area vs. less-massive white dwarfs. This means that massive 
white dwarfs are less luminous for a given temperature, so their evolutionary 
tracks are below those of the less-massive white dwarf stars. 

Theoretical models of the evolution of white dwarfs have been constructed and 
they show that a white dwarf with a mass of 0.6 M 0 will fade to 0.1 L Q in about 
20 million years. Any further reductions in luminosity take progressively longer 


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Temperature (K) 


Figure 3.25. White dwarf evolutionary tracks. 


amounts of time. This means that it will take 300 million years to fade to about 
0.01 L Q) and a billion years to get to 0.001 L Q . It will take about 6 billion years for 
the white dwarf to reach a luminosity of 0.0001 L Q . At this point, the white dwarf 
will have the same temperature and color of the Sun. It will be so faint, however, 
that unless it were within a few pc of the Earth, it would be undetectable. Those 
white dwarf stars with masses greater than 0.6 M Q have more internal heat and 
so will take an even longer time to cool down and grow faint. 

In the case of the Sun, it will eject most of its mass into space and eventually 
end up about the same size as the Earth, but its luminosity will change dramati- 
cally, perhaps only achieving one-tenth the brightness it presently has. As it ages 
further, it will continue to grow even fainter. When about 5 billion years have passed, 
the Sun will only be able to achieve one ten-thousandth of its present luminosity. 
As time passes into the unimaginable future, it will simply fade from view! 

A white dwarf will cool and grow fainter, so it moves downwards and to the 
right on the H-R diagram. The more massive a white dwarf, the smaller and 
fainter it will be. Therefore, the track for a 1 M Q white dwarf will lie below the 
track for less-massive white dwarfs. Note that although a white dwarf may have 
the same temperature as a main-sequence star, it will be fainter because it is 
small, and thus has a smaller surface area. 


3.21 .4 White Dwarf Origins 

It is now believed that most, if not all, white dwarfs have evolved directly from 
the central stars of planetary nebulae. These, in turn, are the former cores of 


Stars 



AGB stars. We saw earlier that during the AGB phase, a star will lose much of 
its mass via a cool stellar wind. If the star strips away sufficient mass to one 
that is lower than the Chandrasekhar limit, a carbon-oxygen rich core of matter 
surrounded by a very thin layer of helium-rich gas is the result. In some cases, 
there may even be a thin outer layer of hydrogen-rich gas. The star and expelled 
gas are now a planetary nebula, and at the moment nuclear fusion ends, a white 
dwarf is born. 43 But even though theory matches well with observations, there 
is still uncertainty as to the mass that the star may have originally been to lose 
enough mass to become a white dwarf star. Current ideas suggest a mass limit of 
8M 0 . Those main-sequence stars that have a mass between 2 and 8M 0 produce 
white dwarfs of mass 0.7 and 1.4M 0 , whereas main-sequence stars less than 
2M 0 produce white dwarfs of mass 0.6 to 07 M 0 . If a white dwarf star has a 
mass less than 0.6 M 0 , the progenitor main-sequence star will have a mass less 
than 1M 0 . What is incredible about these lower-mass stars is that their main- 
sequence lifetimes are so incredibly long, the universe is not yet old enough 
for them to have evolved into white dwarf stars. This means that there are no 
white dwarf stars with a mass less than about 0.6 M Q . The timescale for the 
evolution from giant star to white dwarf can take between 10,000 and 100,000 
years. 

Due to their faintness and small size, white dwarf stars present a challenge to 
observers. There are, of course, many of them in the night sky, and those amateur 
astronomers with large telescopes of, say, aperture 25 cm in diameter and larger 
will have no problem locating and observing them. On the other hand, there are 
a handful that, given the right conditions, can be seen with much more modest 
instruments. These are the ones I shall outline below. The symbol ® indicates 
the size of the white dwarf star as compared with the Earth (thus, 0.5 ffi would 
mean that it is half the size of the Earth). 


The companion star to the brightest star in the sky ( Sirius ) is a white dwarf 
known as the Pup, the first ever to be discovered. It is a difficult, though not 
impossible, star to observe for two main reasons. First, it is overcome by the 
dazzling primary star, and so the light from Sirius often needs to be blocked out 
by some means. In fact, if Sirius B were not a companion to Sirius, it would be 
easily visible in binoculars. Second, its orbit changes over a period of 50 years. 
This means that at certain times it will be too close to Sirius to be detected with 
amateur instruments. The next time it will be at maximum separation is the year 


3.2 1 .5 Bright White Dwarf Stars 


Sirius B 
8.4m 


06'' 45.1'" 
11.2M 0.92® 


— 16° 43' Dec— Jan— Feb 
27,000 K Canis Majoris 


2025. 


Procyon B a Canis Minoris 07 h 39.3 m +05° 13' Dec— Jan— Feb 
10.9m 13. 2M 1-05® 8700 K Canis Minoris 


This dwarf star is not easily visible in small amateur telescopes, having a 
magnitude of 10.8 and a mean separation of 5 arcsecs. Note that it has a low 


1 38 Astrophysics is Easy 

temperature compared with other white dwarfs. It is the second-closest white 
dwarf to the Earth. A challenge to observers. 

o 2 Eradani 40 Eridani B 04 , '15.2 m -07°39' Oct-Nov-Dec 
9.5m 11. 0M 1-480 14,000 K Eridanus 

Even though this is a challenge to split with binoculars, it is nevertheless the 
easiest white dwarf star to observe. The star will be in a prime observing position 
relative to its brighter primary star for the next 50 years or so. What makes 
this system so interesting is that the secondary is the brightest white dwarf star 
visible from Earth. In addition, under high magnification, the white dwarf will 
be seen to have a companion star of its own — a red dwarf star! All in all, a nice 
triple star system. 

Van Maanen's Star Wolf 28 00 ,! 49.1 m +05°25' Aug-Sep-Oct 

12.3m 14. 1M O.9(?) 0 6000 K Pisces 

One of the few stars visible to amateurs, it is a close white dwarf, at only 
13.8 l.y. distant. It is located about 2° south of S (delta) Piscium. Discovered 
by A. Van Maanen in 1917 due to its large proper motion of 2.98 arcsecs per 
year. 


3.22 High-Mass Stars 
and Nuclear Burning 


We now turn our attention to high-mass stars. As you have probably surmised 
by now, the death throes of these stars are very different than and spectacular 
compared with those of low-mass stars. 

Throughout the entire life of a low-mass star (that is, one that is less than 
4M 0 ), only two nuclear reactions occur: hydrogen-burning and helium-burning, 
and the only elements besides hydrogen and helium that are formed are carbon 
and oxygen. Stars that have a zero-age mass greater than 4M 0 begin their lives 
in a similar manner, but theory predicts that due to the increased mass, and 
therefore higher temperatures involved, other nuclear reactions will occur. The 
tremendous crush of gravity is so overwhelming that degeneracy pressure is 
never allowed to come into play. The carbon-oxygen core is more massive than 
the Chandrasekhar limit of 1.4 M 0 , and so the degenerate pressure cannot stop 
the core from contracting and heating. 

The nuclear reactions that take place in the star’s final phase of its life are 
very complex, with many different reactions occurring simultaneously. But the 
simplest sequence of fusion involves what is termed helium capture; this is the 
fusing of helium into progressively heavier elements. 44 The core continues to 
collapse with an accompanying rise in temperature to about 600 million K. At this 
high temperature, the helium capture can give rise to carbon-burning, and the 
carbon can be fused into heavier elements. The elements oxygen, neon, sodium, 
and magnesium are produced. The carbon fusion provides a new source of energy 



Stars 


139 


that, albeit temporarily, restores the balance between pressure and gravity. If the 
star, however, has a mass greater than 8M 0 , even further reactions can occur. In 
this phase, the carbon-burning may only last a few hundred years. As the core 
contracts further, the core temperature reaches 1 billion K, and neon-burning 
begins. In this manner, the neon produced by the earlier carbon-burning reaction 
is used up, but at the same time, there is an increase in the amount of oxygen, 
and magnesium in the star’s core. This reaction lasts as little as 1 year. As you can 
imagine, with each stage of element burning, higher temperatures are reached, 
and further reactions occur; oxygen-burning will occur when the reaches 1.5 
billion K, with the production of sulphur. Silicon-burning can also occur if the 
core reaches the staggering temperature of 2.7 billion K. This reaction produces 
several nuclei, from sulphur to iron. 

Despite the very dramatic events that are occurring inside the high-mass star, 
its outward appearance changes only slowly. When each stage of core nuclear 
fusion stops, the surrounding shell-burning intensifies and therefore inflates the 
star’s outer layers. Then, each time the core flares up again and begins further 
reactions, the outer layers may contract slightly. This results in the evolutionary 
track of the star zigzagging across the top of the H-R diagram. 

Some of the reactions that occur also release neutrons, which are particles 
similar to protons, but they do not have an electric charge. This neutrality means 
that they can, and do, collide with positively charged nuclei and combine with 
them. The absorption of neutrons by nuclei is termed neutron capture. In this 
way, many elements and isotopes that are not produced directly in the fusion 
reactions are produced. 

Each stage during this phase of a high-mass star’s life helps to initiate the 
subsequent phase. As each phase ends due to the star’s using up the specific 
fuel in its core, gravity will cause the core to contract to an ever-higher density 
and temperature, which in turn is responsible for starting the next phase of 
nuclear-burning. In effect, you can think of each stage burning the “ash” of the 
previous one. 

An interesting point to make here is that we tend to think of astronomical 
events taking place over many millions of years. However, theoretical calculations 
have shown that when we are dealing with high-mass stars, events can proceed 
at a very fast pace, with each successive stage of nuclear-burning proceeding at 
an ever-increasing rate. One calculation has been made in detail for 20-25 M 0 
zero-age stars, and the results are very surprising. The carbon-burning stage can 
last for about 600 years, while the neon-burning stage can be as short as 1 year. 
Then things start to speed up! The oxygen-burning lasts only 6 months, and the 
silicon-burning only 1 day! 

At each phase of core-burning, a new shell of material is formed around the 
core of the high-mass star, and after several such stages, say, a very massive 
star of mass 20-25 M 0 , the internal structure of the star can resemble an onion 
(shown in Figure 3.26). 

Nuclear reactions are taking place in several different shells simultaneously, 
and the energy released does so at such a rapid rate that the star’s outer layers 
can expand to an immense size. The star can now be called a supergiant star. 
The luminosity and temperature of such stars are much higher than those of a 
mere giant star. 


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Astrophysics is Easy 



Nonburning hydrogen 
/ Hydrogen fusion 


Helium fusion 
^ Carbon fusion 


Carbon fusion 


Oxygen fusion 


x Neon fusion 
Silicon fusion 


Inert iron core 


Figure 3.26. The multiple-layer structure of an old high-mass star. 


Many of the brightest stars in the night sky are supergiants. These include 
Rigel and Betelgeuse in Orion and Arcturus in Scorpius. Rigel has a temperature of 
1 1,000 K, while Betelgeuse is only 3700 K (or even cooler) and is an example of a 
"red supergiant.” Thus, although Betelgeuse is cooler, it must be correspondingly 
larger for it to be as bright as Rigel. Oddly enough, red supergiants are rare, 
perhaps even rarer than the O-type stars. One current estimate predicts that 
there is only one red supergiant star for every million stars in the Milky Way, 
and only about 200 have ever been studied. 

What makes these stars stand out is their immense size. The radius of Betel- 
geuse has been measured to be about 700 times that of the Sun, or 3.6 astro- 
nomical units. This can be better appreciated this way: if it were placed in the 
Solar System, it would extend past the asteroid belt to about half-way between 
the orbits of Mars and Jupiter. Antares would extend nearly to Jupiter! Alpha 
Herculis is only 2 astronomical units in radius. The record, however, must go 
to VV Cephei, which is an eclipsing binary star. Its radius is a staggering 1900 
times that of the Sun, or 8.8 astronomical units. This means that it would nearly 
extend to Saturn. 

As I mentioned earlier, supergiant stars are quite rare, but fortunately for 
observers, there are some we can see with the naked eye. These are listed below. 


Also known as 119 Tauri, this star has a radius of 2.9 astronomical units and lies 
about 2000 l.y. from us. It has the odd distinction of being classified as both a semi- 
regular and irregular variable star, meaning it is an erratic variable star, and so its 
period is difficult to predict with any certainty. It lies within a field of stars of similar 
brightness, which makes it difficult to locate unless a good star atlas is handy. 

Mu Geminorum HD 44478 06' , 12.3 m +22°54' Nov-Dec-Jan 

6.51m — 4.09M M2 lab Gemini 


3.22.1 Bright Supergiant Stars 


CE Tauri HD 36389 05 ft 32.2 m 

4.38m — 6M M2 lab 


+ 18°35' Nov— Dec— Jan 
Taurus 



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141 


Part of the Gem OBI stellar association, is at the limit of naked-eye visibility as 
observed from an urban location. It lies at a distance of 4900 l.y. 

Other stars that are supergiants and were mentioned in earlier sections are: 
Rigel, Betelgeuse, Antares, Mu Cephei, Eta Persei, i|< _1 Aurigae, and VV Cephei. 

Before we leave supergiant stars, I should mention a class of stars that are 
similar to supergiants, and these are the Wolf-Rayet stars. These are very hot, 
very luminous supergiant stars, similar to O-type stars, but they have very strange 
spectra that show only emission lines and, strangely enough, no hydrogen lines. 
Wolf-Rayet stars are believed to be precursors to the formation of planetary 
nebulae. They are few and far between, with perhaps only 1000 in our Galaxy. 
They have a terrific mass loss, and images from large telescopes show these stars 
surrounded by rich clouds of ejected material. Fortunately for us, there is a very 
bright example that can be easily observed. 

y 2 Vel HD 68273 08 , '09.5 m -47°20' Dec-Jan-Feb 

1.99 v m 0.05 M WC 8 Vela 

The brightest and closest of all Wolf-Rayet stars, y 2 Vel is an easy double, colors 
white and greenish-white. 

This aspect of a supergiant’s life, whereby several layers of nuclear-burning 
occur, resembling the layers of an onion, cannot go on forever, as there is only a 
finite amount of material to burn. Thus, a point comes when the high-mass star 
undergoes yet another change, but this time, with catastrophic consequences. It 
is star death, but in a spectacular manner — a supernova. 


3.23 Iron, Supernovae, and the 
Formation of the Elements 


When nuclei in the core collide and fuse, energy is emitted, and it is this energy 
flowing from the core and surrounding shells of nuclear-burning that supports the 
tremendous weight of material making up a star. The energy is a consequence of 
the strong nuclear force of attraction between neutrons and protons, or nucleons, 
as they are sometimes called. But you may recall that protons also repel each 
other by what is called the weak electric force. This has profound consequences 
for the life of the high-mass star. 

Up to this point, energy has been released (i.e., the energy has been output, but, 
due to the repulsive effect, if any protons are added to nuclei larger than iron, 
which has itself 26 protons, then energy must be input to the system). What this 
means is that any nuclei greater than iron will not release any energy. Therefore, 
the various stages of nuclear-burning end with the production of silicon. After 
that, iron can be formed, but there will be no release of energy associated with 
its formation. The result is an iron-rich core that has no nuclear reactions taking 
place within it. 

Of course, surrounding this inert core of iron will be the various shells of 
nuclear burning. 45 However, this is a state of affairs that cannot continue for 
much longer. 




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Astronomers use a variety of techniques to find out about the life of a star. 
Observations are made, and then theoretical models are devised so that they fit 
the observations. In the case of supernovae, it can be said that most, if not all, 
of what we know about supernovae comes from theoretical and mathematical 
calculations. After all, it is not easy to see what is happening in the central regions 
of a star! You will also see that we are now talking about densities, pressures, 
and velocities that will stagger our comprehension. With this in mind, note 
that the following descriptions of the events in a high-mass star are theoretical 
predictions, albeit ones that seem to fit the observations. 46 

During these final days of a star, the core of inert iron, in which there are no 
nuclear reactions taking place, is surrounded by shells of silicon, oxygen, neon, 
carbon, helium, and hydrogen. The core, which can be thought of as essentially a 
white dwarf star surrounded by the outer layers of a red giant star, is supported 
by the pressure of its degenerate electrons. There is a limit to the mass of a white 
dwarf star — the Chandrasekhar limit — and so when the core surpasses this limit, 
its weight becomes too great to be supported by the degenerate electrons and it 
collapses. 

A consequence of the core contraction is an increase in the density, which 
in turn gives rise to a process called neutronization. This is a process where-by 
electrons react with the protons in the iron nuclei to form neutrons. This is 
shown below. 


e +p->n+v 

Each neutronization reaction will also produce a neutrino. Now more and more 
electrons will react with the protons, and so there are fewer left to support 
the core, and so resist the compression. This results in a speeding up of the 
contraction and actually could be better termed a collapse of the core. It only 
takes about a second for the core to collapse from a radius of thousands of 
kilometers to about 50 km. Then, in only a few seconds, it shrinks down to a 
5-kilometer radius. The core temperature also increases during this time to about 
500 million K. The gravitational energy released as a result of the core collapse is 
equal to the Sun’s luminosity for several billion years. Most of this energy is in 
the form of neutrinos, but some is also in the form of gamma rays, which are 
created due to the extremely hot core temperature. These gamma ray photons 
in turn have so much energy that when they collide with the iron nuclei, the 
nuclei are broken down into alpha particles (which are 4 He nuclei). This process 
is called photo-disintegration. 

After a short interval of time, which is thought to be about 0.25 seconds, the 
central 0.6 to O.8M 0 of the collapsing core will reach a density equal to that of 
atomic nuclei, that is, some 4 x 10 27 kg m“ 3 . At this point, the neutrons become 
degenerate and strongly resist any further attempts of compression. To get an 
idea of what this density means, the Earth would have to be compressed to a 
sphere 300 meters in diameter. For all intents and purposes, the core of the star 
can now be thought of as a neutron star, and the innermost part of the core 
suddenly becomes rigid, and the contraction abruptly halts. This innermost part 
actually rebounds outward and pushes back against the rest of the infalling core, 
driving it outward in a pressure wave. This is called the core bounce, and it is 
illustrated in Figure 3.27. 


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143 


T 









\ 


Step 1: the iron core collapses Step 2: the neutron-rich 

core rebounds 


Step 3: the shock wave 
moves outward 


Figure 3.27. Evolution of a supernova explosion. 


The core also cools at this stage, and this causes the pressure to decrease 
significantly in those regions that surround the core. If you recall, there is a 
balancing act between the pressure pushing upward and gravity falling downward, 
and a consequence of this reduced pressure is that the material surrounding 
the core now falls inward at a velocity close to 15% of the speed of light. This 
inward-moving material encounters the outward-moving pressure wave, which, 
incidentally, can be moving at one-sixth the speed of light. In just a fraction of a 
second, the falling material now moves back outward toward the star’s surface. 

Surprisingly enough, this wave of pressure would soon die out, long before it 
reached the surface of the star, if it were not for the fact that helping it along 
is the immense number of neutrinos that are trying to escape the star’s core. 
The upward-moving wave of pressure speeds up as it encounters the less-dense 
regions of the star and achieves a speed in excess of the speed of sound in the 
star’s outer regions. The pressure wave now behaves like a shock wave. 

These neutrinos actually escape from the star in a few seconds, but it takes 
a few hours for the shock wave to reach the surface. Most of the material of 
the star is pushed outward by this shock wave and is expelled from the star at 
many thousands of kilometers per second. The energy released during this event 
is a staggering 10 46 J, which is 100 times more than the entire output of the Sun 
during the last 4.6 billion years. It will surprise you to know that the visible light 
we observe is only about 1% of the total energy released during the event. 

Recent studies have proposed that up to 96% of the material making up the 
star may be ejected into the interstellar medium that, of course, will be used 
in future generations of star formation. But before this matter is ejected, it is 
compressed to such a degree that new nuclear reactions can occur within it, 
and it is these reactions that form all the elements that are heavier than iron. 
Elements such as tin, zinc, gold, mercury, lead, and uranium, to name a few, are 
produced, and this has profound implications because it means that the stuff 
that makes up the Solar System, the Earth, and in fact us, was formed long ago 
in a supernova. 



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The expansion of the star’s surface due to the shock wave is the cause of the 
tremendous increase in luminosity that we observe, but after several months, 
the surface will cool, and so the brightness will fade. During this later-stage 
time, the main source of the supernova’s light is in fact the radioactive decay of 
nickel and cobalt nuclei, which were produced in the supernovae event. These 
decaying nuclei are able to keep the supernova shining for many years. 

As an observer, it would be delightful if you could go out and just pick a 
supernovae a that you wish to look at. Life is not like that, however from a 
statistical point of view, there should be about 100 supernovae a year in our 
Galaxy, and so you would think that you would have a good chance of observing 
one. The most recent bright supernova, in 1987, was in fact in another galaxy 
completely — The Large Magellanic Cloud — and the last bright one in our Galaxy 
was seen several hundred years ago. So why is this? The answer is simple. As we 
have seen in earlier sections, our Galaxy is filled with dust and gas, and it is this 
material that blocks out the light from any supernova that may be occurring. 
That is not to say that we will never see a supernova — far from it; but we cannot 
predict with any certainty when one will occur, although there are a few stars 
that we should keep an eye on: Betelgeuse and Eta Carinae. 

We can, however, see the remains of a supernova, the supernova remnant. 


The supernova remnant (usually abbreviated to SNR) represents the debris of 
the explosion, the layers of the star that have been hurled into space, and the 
remains of the core that will now be a neutron star. The visibility of the remnant 
actually depends on several factors: its age, whether there is an energy source to 
continue making it shine, and the original type of supernova explosion. 

As the remnant ages, its velocity will decrease, usually from 10,000 km s -1 to 
maybe 200 km s -1 . It will, of course, fade during this time. A few SNRs have a 
neutron star at their center that provides a replenishing source of energy to the 
far-flung material. The classic archetypal SNR that undergoes this process is the 
Crab Nebula, Ml in Taurus. What we see is the radiation produced by electrons 
travelling at velocities near the speed of light as they circle around magnetic 
fields. This radiation is called Synchrotron radiation and is the pearly, faint glow 
we observe. Some SNRs glow as the speeding material impacts dust grains and 
atoms in interstellar space, while others emit radiation as a consequence of the 
tremendous kinetic energies of the exploding star material. 


Also known as the Veil Nebula (Western Section). This is the western portion of 
the Great Cygnus Loop, which is the remnant of a supernova that occurred about 
30,000 years ago. It is easy to locate because it is close to the star 52 Cygni, though 
the glare from this star makes it difficult to see. The nebulosity we observe is 


3.23.1 Supernova Remnants 


Caldwell 34 NGC 6960 20 h 45.7 m 

•3-5 © 70 1 6' 


•3-5 


+30° 43' Jul— Aug— Sep 
Cygnus 


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145 


the result of the shockwave from the supernova explosion’s impacting on the 
much-denser interstellar medium. The actual remains of the star have not been 
detected. 


Caldwell 33 NGC 6992 20 h 56.4 m +31°43' Jul-Aug-Sep 

•2-5 0 6O|8' Cygnus 

Also known as the Veil Nebula ( Eastern Section). A spectacular object when 
viewed under good conditions. It is the only part of the loop that can be seen in 
binoculars and has also also been described as looking like a fish hook. Using a 
telescope, it becomes apparent why the nebula has been named the Filamentary 
Nebula, as lacy, delicate strands will be seen. However, there is a downside: it is 
notoriously difficult to find. Patience, clear skies, and a good star atlas will help. 
A showpiece of the summer sky (when you have finally found it). 

IC 2118 05 h 06.9 m -07°13' Nov-Dec-Jan 

•3-5 0 180 1 60' Eridanus 

Also known as the Witch Head Nebula. This is a very faint patch of nebulosity, 
which is apparently the last of a very old supernova remnant. It resembles a 
long ribbon of material, which can be glimpsed with binoculars. It is glowing 
by reflecting the light of nearby Rigel. Very rarely mentioned in observing 
guides. 


Messier 1 NGC 1952 05 h 34.5 m +22°01' Nov-Dec-Jan 

• 1-5 0 6|4' Taurus 

Also known as the Crab Nebula. The most famous supernova remnant in the 
sky, it can be glimpsed in binoculars as an oval light of plain appearance. 
With telescopes of aperture 20 cm, it becomes a ghostly patch of grey light. In 
1968, in its center was discovered the Crab Pulsar, the source of the energy 
responsible for the pearly glow observed, a rapidly rotating neutron star that 
has also been optically detected. The Crab Nebula is a type of supernova 
remnant called a plerion, which, however, is far from common among supernova 
remnants. 

Sharpless 2-276 05 h 56.0 m — 02°00'y Nov— Dec— Jan 

•6.5 ? 0 600' Orion 

Also known as Barnard’s Loop. It seems to be the remains of a very old supernova. 
Often mentioned in books, but very rarely observed, this is a huge arcing loop of 
gas located to the east of the constellation Orion. It encloses both the sword and 
belt of Orion, and if it were a complete circle, it would be about 10° in diameter. 
The eastern part of the loop is well defined, but the western part is exceedingly 
difficult to locate and has never to my knowledge been seen visually. Needless 
to say, perfect conditions and very dark skies will greatly heighten the chances 
of its being seen. 


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Astrophysics is Easy 


3.23.2 Supernova Types 

Before we bid a fond farewell to Supernovae, and say hello to the final phase of a 
star’s life, I should mention (albeit briefly) that there are two types of supernovae. 
The classification system used to distinguish the two types is a rather obtuse (for 
the non-professional astronomer) system based upon whether the supernova has 
emission lines of hydrogen in its spectrum. Type I supernovae do not have these 
lines, whereas Type II do. The supernova we have previously discussed is a Type 
II supernova. This class of supernova involves the final death throes of a massive 
star. These stars, as we have seen, have quite a lot of hydrogen left in their outer 
layers, hence the classification as a Type II. Type I stars, on the other hand, do 
not have hydrogen emission lines and can be further divided into Types la, lb, 
and Ic. Type la has absorption lines of ionized silicon, and Types lb and Ic do 
not. The difference between lb and Ic is also a spectroscopic one, in that the 
former has a helium absorption line, whereas the latter does not. 

There is a further twist to this story. Types lb, Ic, and II are massive stars, but 
Type I stars have had their outer layers stripped away either by a strong stellar 
wind or by the action of a nearby star (so that the progenitor supernova star is 
in fact part of a binary star system), and mass is transferred from one star to 
another. Furthermore, all three are usually found near sites of star formation. 
This is to be expected, as we know that massive stars have short lives, and so we 
do not expect them to move far from their birthplace. 

Type I supernovae are a different beast altogether. They are usually (but 
not always) found in galaxies where star formation may be minimal or has 
even stopped altogether. You can now see that this implies that they originate 
not from the final phases of massive stars, as described above, but from some 
other, new phenomenon. Actually, in an earlier section, we discussed stars that 
are the originators of Type la supernovae. These are white dwarf stars that 
literally explode by thermonuclear reactions. Now, you may think that this is 
contradictory to what you read earlier, where I said that white dwarf stars do not 
have any nuclear reactions occurring within them. This is absolutely correct, but 
in this case, the carbon-rich white dwarf star is part of a binary system where 
the other star is a red giant. 

Recall that as a red giant star evolves, its outer layers expand and it can 
overflow what is called its Roche Lobe. This is the region around a star in 
which the gravity of the star dominates. Any matter within the Roche Lobe is 
gravitationally bound to the star, but if the Roche Lobe is filled, then matter 
can overcome the gravitational attraction of the star and in fact “flow” or be 
transferred to a companion star . 47 This material then falls onto the white dwarf 
star. A consequence of this extra material is that the Chandrasekhar limit can be 
reached, and the increase in pressure will cause carbon-burning to commence 
deep in the star’s interior. This is, of course, accompanied by a resulting increase 
in temperature. 

Normally, this temperature increase would mean a further increase in pressure, 
resulting in an expansion of the white dwarf, resulting in the star’s cooling 
down, and the carbon burnings ceasing. However, as we have seen, white dwarfs 
are not normal. They are made of degenerate matter, which means that the 


Stars 


increase of temperature results in the carbon-burning reactions proceeding at 
an ever-increasing rate. This is reminiscent of the helium flash process seen in 
low-mass stars mentioned earlier. The temperature soon gets so high that the 
electrons in the white dwarf become non-degenerate, and, simply put, the white 
dwarf blows itself to bits! 

So, we can see that Type I supernovae involve nuclear energy and emit more 
energy in the form of electromagnetic radiation, 48 whereas Type II involve gravi- 
tational energy and emit an enormous number of neutrinos. 

All of this, of course, has no bearing on the amateur astronomer who wishes 
to observe a supernova, as this will not help or hinder you. Nevertheless, it is 
important to know in case you read in journals or on the WWW that the latest 
supernovae is of Type lb (or Ic or la or even Type II). 

We now move on to the final phase of a star’s life, the end result of millions, 
and even billions, of years of stellar evolution, the end of a long journey. 


3.24 The End Result of High-Mass 
Stars' Evolution: Pulsars, 
Neutron Stars, and Black 
Holes 


3.24.1 Pulsars and Neutron Stars 

The endpoint of a star’s life is now in sight, and although these objects are 
probably forever beyond the vision of amateur astronomers, it is important that 
we discuss them for the sake of completeness. These objects are either very small 
(maybe 10 km in radius) or invisible, so for all intents and purposes they are 
not observable with amateur equipment. 49 They represent the conclusion of a 
star’s evolution and until fairly recently, had never been seen, only predicted. 
The fascinating properties of these objects could fill a book in itself, so I shall 
just briefly describe those properties that are relevant to the evolutionary story. 

Recall that in a Type II supernova the central O.6M 0 of the collapsing core 
has a density equal to that of the nuclei of atoms, and the neutrons become 
degenerate. This central core region has become a neutron star. In fact, after a 
supernova explosion has flung all the outer layers of the star into space, what 
remains (usually) is just the central core region. These neutron stars were actually 
predicted as far back as 1939 by Robert Oppenheimer and George Volkoff, who 
calculated the properties of a star made entirely of neutrons. 

The actual structure of the star is not completely known, but there are many 
theoretical models that accurately describe the observations. Many of their 
properties are similar to those of white dwarf stars. For instance, an increase 
in the mass of a neutron star will result in a decrease in radius, with a range 
of radii from 10-15 km. The mass of a neutron star can be from 1.5-2.7M 0 . 


1 48 Astrophysics is Easy 

But, of course, these figures depend on the calculation being used at the time. 
Nevertheless, they give a good picture of the star. 

Two properties of neutron stars that we can describe in confidence are its 
rotation and magnetic field. A neutron star rotates at a very rapid rate, as many 
as hundreds and even thousands of times per second. It does this because of a 
law of physics called the conservation of angular momentum. Although it is a 
complicated law, it is easy to visualize. Just picture an ice skater spinning around; 
as she pulls in with her arms, she spins faster. It is the same with the neutron 
star. The Sun rotates about once every 30 days, but if it were shrunk to the size 
of a neutron star, it would have a rotation rate of 1000 times per second. We also 
know that every star has a magnetic field, but imagine compressing that field to 
the size of a neutron star with the result that it would be enormous. Again, using 
the Sun as an example, its magnetic field would increase by 10 billion times if it 
shrunk to the size of a neutron star. The strength of the field of a typical star is 
about 1 tesla, whereas in a neutron star it can be as high as 100 million tesla. 

Some neutron stars are believed to be in a binary system, and material can 
be transferred from the companion star onto the magnetic pole regions of the 
neutron star with the matter travelling at perhaps nearly half the speed of light. 
The material literally crashes onto the star and results in “hot spots.” The temper- 
ature of these hot spots is high, in the range 10 8 K, and can result in the emission 
of X-rays. In fact, to casually say they emit X-rays is misleading, as the amount 
of X-ray emission is tremendous. The total amount of X-ray luminosity can be as 
high as 10 31 watts, nearly 100,000 times greater than the total amount of energy 
emitted by the Sun at ALL wavelengths! These x-ray bursters typically flare up 
and last from a few hours to a few days. Each burst lasts only a few seconds, but 
then declines in energy and brightness. This type of binary system is called an 
X-ray binary pulsar, and examples are Hercules X-l and Centaurus X-3. 

This leads rather nicely to the subject of pulsars. In 1967, a young graduate 
student at Cambridge (Jocelyn Bell) discovered a source of very evenly spaced 
pulses of radio emission. The period of the pulses was 1.337 seconds and was 
very constant to an accuracy of about 1 part in 10 million. The object, designated 
PSR 1919 + 21, was the first pulsar discovered! The problem was trying to explain 
what this object was. Some theories predicted a neutron star that pulsated in a 
manner similar to that of a Cepheid variable star, where its size actually changed. 
One proposal even suggested that these pulsars were in fact messages from an 
alien civilization. Not surprisingly, this last idea was discounted. Another model 
was that of a rotating white dwarf star. All of these plausible explanations were 
eventually discounted and the correct one accepted. 

The generally accepted model of a pulsar is one in which the magnetic axis 
of the neutron star is tipped with respect to its rotation axis. Very energetic 
particles travel along the magnetic field lines and are literally beamed out from the 
magnetic poles. As the neutron star rotates around its rotation axis, the beamed 
radiation sweeps across the Earth and the pulse detected. In some instances, two 
pulses can be observed per rotation if the beams from both magnetic poles sweep 
past the Earth. As time passes, the period of a pulsar increases. For instance, 
a pulsar with a period of 1 second will slow down to a rate of 2 seconds in 
30 million years. One observational point to make is that although we know of 
many pulsars, there are none that have periods of, say, 5 seconds or longer; this 


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149 


would imply that the pulse mechanism must cease after a period of time. So, 
neutron stars only exist as pulsars for the first tens of millions of years after the 
supernova explosion. 

I mentioned earlier that neutron stars are the remains from a supernova 
explosion, and that pulsars are rotating neutron stars. So, we would expect to 
find pulsars at the center of supernova remnants, or SNRs. We do, but so far 
there are only 3 known SNRs with associated pulsars. There are two reasons for 
this. To detect a pulsar, the beams have to sweep past the Earth, and if they 
do not we will not detect them. Second, the supernova remnant will only last a 
relatively short time, perhaps 100,000 years, before it merges into the interstellar 
medium and disappears from view. On the other hand, a pulsar can last for 
millions of years. So, many of the pulsars we observe now are old, with their 
SNRs having dispersed. 

An example of a pulsar at the center of an SNR, and probably the most famous, 
is the one in the Crab Nebula designated PSR 0531-21. In fact, the energy from 
the pulsar is responsible for the pearly glow and appearance of the nebula, and it 
is caused by synchrotron radiation produced by high-velocity electrons spiralling 
around the magnetic field. An SNR that has a filled-in appearance as opposed to 
a shell-like appearance is termed a plerion. 

I mentioned in the section on white dwarfs that there is a limit to a white 
dwarfs mass (the Chandrasekhar limit), beyond which the star cannot support 
the weight of the material making it up. Not surprisingly, there is also a limit to 
the mass a neutron star can support. Current estimates put this figure at about 
2-3 M q . In some supernovae, the most massive outer layers may not have been 
dispersed into space during the explosion, and matter may fall back onto the 
already-dense core. This extra material may push the neutron star core above its 
limit, and neutron degeneracy pressure will not be able to fend off gravity. 

The core will continue to collapse catastrophically, and not even the increasing 
temperature and pressure can halt the inevitable result. In fact, according to 
Einstein’s famous equation, E = me 2 , energy is equivalent to mass and so the 
energy associated with the incredible pressure and temperatures concentrated in 
the now-tiny core acts like additional mass, thus hastening the collapse. To the 
best of our knowledge, nothing can stop the crush of gravity. The core collapses 
without end, forming a black hole. We have reached the end of a star’s life. 


3.24.2 Black Holes 

We now discuss an object that everyone has heard about, but not many really 
understand — a black hole. It is one of those things that has gripped the public’s 
imagination, what with its fabled inescapable pull of gravity and its possible use 
as a means of stellar transportation. However, it will surely come as a surprise 
to many of you to know that although a rigorous description of a black hole 
would entail a thorough background in Tensor Calculus and General Relativity, 
the broad description of such an exotic object is quite simple, and it is very easy 
to calculate a few of its basic characteristics. Let us begin. 

Following the previous section’s description of the formation of a neutron 
star and subsequent supernova, if the core of the star contains about 3 or more 


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Astrophysics is Easy 


solar masses, nothing will stop its collapse even beyond the neutron degeneracy 
stage. In fact, the core will collapse to an object with zero radius! Consider this 
statement for a moment. Something with zero radius has no physical size; it is 
not very tiny, or even really, really tiny. It has no size at all. 

The core collapses, and thus its density, and thus surface gravity increase. If 
it collapses to zero size, then the gravity becomes infinite, and that is a lot of 
gravity. This entity, which has no physical size yet has infinite gravity, is called 
a singularity. 

Before we go any further, we need to make a slight detour to discuss escape 
velocity. This is the velocity an object needs to escape the pull of gravity of 
a celestial object. For instance, a spaceship has to go about 11 kilometers per 
second to escape from the Earth. What dictates the value for the escape velocity 
depends on two things: the mass of the celestial body and the distance from the 
object escaping to the center of the celestial body. Thus, if something were very 
massive, or very small, then the escape velocity would increase. A point may be 
reached when even light, which is the fastest-moving thing in the universe, may 
not be able to achieve escape velocity, and since the light would never escape, 
the object would never be seen. 50 

Let us return now to the singularity. There will be an area of space (usually 
spherical) surrounding the singularity where the escape velocity will be so high, 
light cannot escape. This sphere of space within which the escape velocity is 
equal to, or higher than, the escape velocity of light is what we call a black hole. 51 

Thus, inside a black hole, you would have to go faster than light to escape its 
gravitational pull; outside a black hole, you would not. The boundary between 
these two regions is called the event horizon; any event that occurs within the 
horizon is forever invisible to an outside observer. 

To fully understand how and why black holes exist as I said earlier is beyond the 
scope of this book, but we should mention that the great scientist Albert Einstein 
and his theory of General Relativity really started the whole ball rolling. He was 
the first person to combine space and time into a single entity — space-time — and 
it is his equations that showed that gravity can be portrayed as a curvature of 
space-time. Astronomer Karl Schwarzschild used and solved Einstein’s equations 
to give the first-ever general relativistic description of a black hole. We now call 
them Schwarzschild black holes, as his solutions are for non-rotating, electrically 
neutral black holes, to distinguish them from rotating, charged black holes! 

What Schwarzschild showed is that, if the conditions are right (say, if matter is 
packed into a small enough volume of space) then the space-time can curve back 
on itself. This means that an object (or light) can follow a path (also known as a 
geodesic) into a black hole, but inside the black hole, the curvature of space-time 
is so extreme, there exists no path leading out. Once in, you stay in! 

The event horizon is the boundary between the universe and the forever- 
isolated region of extreme curved space-time, 52 also known as a black hole. The 
radius of the black hole, that is, the distance from the singularity to the event 
horizon, is called the Schwarzschild Radius, R s . 

We now have our full description of a black hole: the singularity, the 
Schwarzschild Radius, and the event horizon. 


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One point that I think must be made is that many people believe, erroneously, 
that black holes will just vacuum you up regardless of your distance from them. 
This is wrong. For instance, if the Sun were to suddenly turn into a black hole, 
of radius 3 km, the orbits of the planets would not change. It would get dark 
and very cold, I admit, but the gravitational pull would remain the same. The 
gravitational effect of a black hole only becomes extreme if one gets very close 
to it. Provided we are at a reasonable distance away, there is nothing, relatively 
speaking, to worry about. 53 

There remains, however, one small problem: if black holes are literally invisible 
to us, how can we ever detect them? 

We do this by looking for the effect that they have on objects nearby! First, 
search for a star whose motion, determined by measuring the Doppler shift in 
its spectrum, indicates that it is a member of a binary star. (The Doppler shift is 
the change in the appearance of light from an object that is moving away from 
or toward an observer.) If it proves possible to see both stars, give up on that 
object. The search is for a binary system where one companion is invisible, no 
matter how powerful the telescope used. However, just because it is unseen does 
not mean it is a black hole candidate; it may just be too faint to be seen, or the 
glare from the companion may swamp out its light. It could even be a neutron 
star. 

Thus, further evidence is needed to determine if the invisible companion has a 
mass greater than that allowable for a neutron star. Kepler’s laws are used at this 
point to determine whether the star, or rather the invisible object, has a mass 
greater than 3 solar masses. If this is so, then the unseen companion may be a 
black hole. Further information is still needed, however, and this may appear in 
the form of X-rays, which can arise either from material flowing from one star 
into the black hole or from an accretion disc that has formed around the black 
hole. Either way, the presence of X-ray emission is a good indicator that a black 
hole may be the unseen companion object. 

Of course, the measurements as just stated are a bit more complicated 
than this. For instance, it is known that neutron stars can emit X-rays 
and have an accretion disc. So, careful analysis of the data is necessary. 
However, a few candidates are known, and one is even visible to the amateur 
astronomer. . .or perhaps I should say that the companion star to the black hole is 
visible! 


Box 3.4: The Size of a Black Hole 

To determine the approximate radius of a black hole, known as the Schwarzschild 
radius (R Sch ), we need to know the progenitor mass, in terms of the Sun’s mass, M 0 . 
The radius is given by the very simple formula: 

Rsch ~ 3 

where R Sch is in kilometers. 



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Astrophysics is Easy 


Example: 

Betelgeuse, in the constellation Orion, has a mass of ~ 20 M 0 . Determine the radius 
of a black hole that may form when the star eventually dies. 

Rsch ~ ^ M 0 
Rsch ~ 3 x 20 

~ 60 km 

Thus, a 20 M 0 star could form a black hole of diameter 120 km. 


The point at which Einsteinian 
Gravity takes over from 
Newtonian Gravity 


At a large distance, a black hole exerts a gravitational force according to Newton’s 
Law. However, a point is reached whereby Newton’s Laws are no longer valid, and the 
gravitational effects are now explained using Einstein’s General Relativity. The distance 
from a black hole where this change occurs is approximately 3 R Sch . 

Example: 

Betelgeuse could form a black hole with a Schwarzschild radius of approximately 
60 km. At what distance from the black hole does gravity increase from what Newton’s 
Law predicts? 

Using the above formula, this distance is approximately 3 R Sch , thus: 

~ 3 x 60 
= 180 km 

At 180 km from the black hole, the gravitational force will increase to considerably 
more than that predicted by Newton’s Law. 


Cygnus X— 1 HDE 226868 19 h 58.4 m 35°12' July 21 

This is one of the strongest X-ray sources in the sky and possibly the most 
convincing candidate for a black hole. Its position is coincident with the star 
HDE 226868, which is a BOIb supergiant of magnitude 9. It lies about 0.5° ENE 
of Eta Cygni. It is a very hot star, of around 30,000 K, and analysis shows it is a 
binary with a period of about 5.6 days. Observations by satellite have detected 
variations in the X-ray emission on a time scale of less than 50 milliseconds. The 
estimated mass of the unseen black hole companion is in the range of 6 to 15 
solar masses. This would mean that it has a maximum diameter of about 45 km. 

Other black hole binary systems are: 



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Star 

T yP e 

Orbital Period 
days 

Black Hole Mass 
Estimates M Sun 

LMC X-3 

B main sequence 

1.7 

4 to 11 

V616 Mon 

K main sequence 

7.8 

4 to 15 

V404 Cyg 

K main sequence 

6.5 

> 6 

Nova Sco 1 994 

F main sequence 

2.4 

4 to 15 

Nova Velorum 

M dwarf 

0.29 

4 to 8 


We have finished our brief, but fascinating look at the lives of stars. It is now 
time to look at bigger things — Galaxies! 


Notes 


1. We discuss the proton-proton chain in much greater detail in the sections 
on the Sun and the main sequence. 

2. See the section on the Sun for a full discussion on gravitational equilibrium. 

3. When astronomers refer to a star’s following a specific evolutionary track, 
or moving on an H-R diagram, what they really mean is the star’s luminosity 
and/or temperature changes. Thus, the position of the star on the H-R 
diagram will change. 

4. The theoretical calculations were developed by the Japanese astrophysicist 
C. Hayashi, and the phase a protostar undergoes before it reaches the main 
sequence is called the Hayashi Phase. 

5. There are a few examples of nebulae in which protostars are currently 
forming and which are observable in the section on emission nebulae. You 
will not, however, see protostars, just the region within which they reside. 

6. We shall see in a later section why the Sun is opaque. 

7. Recall from an earlier section that the luminosity is proportional to the 
square of the radius and to the fourth power of the surface temperature. 

8. There is no mass-luminosity relationship for white dwarfs, giant stars, or 
supergiant stars. 

9. This figure of 0.08 M G is about 80 times the mass of Jupiter. 

10. By “ordinary” I mean matter composed of atoms, to distinguish it from 
“dark matter,” whatever that may be! 

11. Stars named after the FU Orionis prototype are also worth observing. It is 
now believed that the activity of FU Orionis (and similar stars) is related to 
the T Tauri variables. T Tauri variations may result from instabilities within 
and interactions with the surrounding accretion disk. FU Orionis activity 
is caused by a dramatic increase in instability due to the dumping of large 
amounts of material onto an accompanying star. Many astronomers believe 
that all T Tauri stars probably go through FU Orionis-type behavior at least 
once in their development. 

12. During 2006, research indicated that some stars do form in isolation. 
However, the consensus outlined above is the one we will adopt. 


1 54 Astrophysics is Easy 

13. Remember that hydrogen-burning is a characteristic of stars on the main 
sequence. 

14. See next section. 

15. Located within Sagittarius are numerous open clusters. Only the brightest 
are listed here. 

16. See Section 3.4.2 in this chapter for a discussion on stellar associations. 

17. We are talking about spiral galaxies here, not elliptical. Elliptical galaxies 
are believed to be the results of mergers between spiral galaxies where the 
rate of star formation is very low. For further details see chapter 4. 

18. It isn’t a cloud at all; this is just the name that ancient astronomers gave 
the galaxy before they knew what it really was! 

19. This is, of course, an exaggeration, as it is only in the last ten years that 
astronomers have solved (possibly!) the problem of the solar neutrino, as 
we shall see in a later section. 

20. There are many excellent books available that are devoted totally to the 
Sun (see appendices for a list of suitable texts). 

21. A plasma is a collection of positively charged ions and free electrons. 

22. In stars that are more massive than the Sun, the fusion of helium occurs 
via a different series of reactions called the CNO cycle. We shall discuss 
this later. 

23. This means that averaged over the distance from the core to the surface, 
a “photon” travels about 0.5 m per hour, or about 20 times slower than a 
snail. 

24. For a very detailed listing of beautiful double and multiple stars, see “Field 
Guide to Deep Sky Objects” by the author. 

25. This is regardless of whether the binary star system is a visual, spectro- 
scopic, eclipsing, or astrometric binary star system. 

26. There are literally thousands of double, triple, and multiple star systems in 
the sky, all within reach of the amateur astronomer. The list that follows 
is a taste of what awaits the observer. If your favorite star is omitted, I 
apologize, but to include them all would be impossible. 

27. Note that the position angle and separation are quoted for epoch 2000.0. 
With double stars that have small periods, these figures will change appre- 
ciably. 

28. The semimajor axis is a term used in elliptical orbits (as nearly all orbits 
are). It is defined as half of the longer axis of an ellipse, and it is the average 
separation of the two stars. 

29. It is this wobble that is used to detect unseen planets around stars! 

30. Note that a A + a B add up to the semimajor axis, a, as used in Kepler’s Law. 

31. We shall discuss this remarkable fact in later sections. 

32. See the next section for a discussion on red giant stars. 

33. Recall that for hydrogen-burning to start, the temperature has to reach 
about 10 million K, whereas for helium -burning, the temperature has to 
achieve a staggering 100 million K. 

34. If you are a particle physicist! 

35. The time for a star to complete one cycle in its brightness variation is called 
its period. Thus, for 8 Cephei, its period is 5.5 days. 


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36. Population I stars are bright supergiant stars, main-sequence stars with 
high luminosity, such as O- and B-type stars, and members of young open 
star clusters. Molecular clouds are often found in the same location as 
Population I stars. They are usually located in the disc of a galaxy and 
concentrated in the spiral arms, following nearly, but not always, circular 
orbits. Population I stars include stars with a range of ages, maybe 10 
billion years old, or one year old. Population II stars, on the other hand, 
are usually old stars. Examples include RR Lyrae stars and the central stars 
of planetary nebulae. This type of star has no correlation with the location 
of the spiral arms. They are also found in globular clusters, which are 
almost entirely in the halo and central bulge of the Galaxy. Therefore, they 
represent the oldest stars, which formed very early in the history of the 
Galaxy. 

37. A list of organizations is in the index. 

38. Current theories predict that low-mass M-type stars will stay on the main 
sequence for trillions of years!!! 

39. In stars that are hotter and have a higher mass than the Sun, the chain of 
reactions that leads to hydrogen fusion is called the CNO cycle, where C, 
N, and O stand for carbon, nitrogen, and oxygen, respectively. The amount 
of energy produced in this reaction is exactly the same as that produced by 
the proton-proton reaction discussed earlier, but it occurs at a much more 
rapid rate. 

40. On one occasion, it remained at minimum magnitude for 10 years! 

41. The name “Planetary Nebulae” was first applied to these objects by 
Herschell, who thought that the nebulae looked like Jupiter when seen in a 
telescope. 

42. See Appendix 1 for a full explanation of degeneracy. 

43. Recent observations have detected a star (V4334 Sagittarii) that was well 
on its way to being a white dwarf when it underwent a final helium flash, 
grew to red-giant size once again, and is now ejecting more gas. Another 
star that has shown similar behavior is V605 Aquilae. 

44. Some helium nuclei do remain in the star’s core, but these are insufficient 
to initiate helium-burning to any great degree. 

45. The entire energy-producing region in the star is now in a volume about 
the same size as the Earth. — one million times smaller in radius than the 
size of the star. 

46. Do not think that astronomers know everything about supernovae. Prior 
to the famous supernovae in 1987 (“SN1987A”), astronomers believed that 
only red supergiants could form supernovae. They were put into some 
confusion when it was discovered that the progenitor star of SN1987 was a 
blue supergiant! 

47. The subject of binary stars, Roche Lobes, and other assorted ephemera 
could fill a book in itself! Any interested readers will find references to 
books on such topics in the appendices. 

48. As there is no core collapse in a Type I supernova, there will be no neutrinos 
emitted. 

49. No doubt I shall soon be corrected on this point, when an amateur images 
the Crab Nebula Pulsar. It is only a matter of time! 


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50. In 1783, the British astronomer Rev. John Mitchell realized that using 
Newton’s Laws of Gravity, a situation could occur whereby an object 500 
times the radius of the Sun but with the same density would have an escape 
velocity greater than the speed of light! Although he didn’t know it, he was 
talking about a black hole. 

51. In some cases, a supernova remnant that does not have a central pulsar or 
neutron star may have a black hole at its center. 

52. There are some relativists who propose that in an unimaginably distant 
future, black holes could “evaporate.” We will be long gone for this issue 
to worry us. 

53. That is, if we ignore the immense amount of radiation being formed around 
a black hole, and the debris from stars that have been literally torn apart. 


CHAPTER FOUR 


Galaxies 


4.1 Introduction 


We now discuss objects that every amateur astronomer has usually seen at least 
a handful of — galaxies. 1 

However, for a majority of amateurs, galaxies tend to remain faint and elusive 
objects, and I would not be far wrong if I said that perhaps only 15 to 20 galaxies 
are ever observed by 99% of amateur astronomers. It often comes as a surprise 
to know that, with the proper optical system and optimum seeing conditions 
(and a copy of this book), there are in fact many more that are within the reach 
of even the smallest telescopes or binoculars. In fact, a few are even visible to 
the naked eye if you know where to look! 

Galaxies are vast, immense collections of stars, gases, and dust. Indeed, they 
are the source of all stars, because stars are not born outside of galaxies. 2 The 
number of stars in galaxies varies considerably; for instance, in some giant 
galaxies, there may be over a trillion (10 12 ) stars — a number that may stagger 
your mind. On the other hand, in small dwarf galaxies, such as Leo I, there may 
be only a few hundred thousand. 


157 




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4.2 Galaxy Types 


Galaxies come in a variety of shapes and sizes, but the vast majority can be 
grouped into a few distinct classifications. When astronomers first began studying 
galaxies, the most obvious characteristic that immediately became apparent was 
their shape, or morphology. Broadly speaking, galaxies can be classified into 
three major categories: 

Spiral galaxies appear as flat white discs with yellowish bulges at their centers. 
The disc regions are occupied by dust and cool gas, interspersed with hotter 
ionized gas, as is the case in the Milky Way. Their most obvious characteristic is 
the beautiful spiral arms. 

Elliptical galaxies are somewhat redder and more rounded in appearance like 
a football. 3 Compared with spiral galaxies, ellipticals contain far less cool gas and 
dust, but very much more hot ionized gas. 

Galaxies that appear neither disc-like nor rounded are classified as irregular 
galaxies. 

Some spiral galaxies exhibit a straight bar of stars that cuts across the center 
with spiral arms curling away from the ends of the bars. Galaxies with these 
features are known as barred spiral galaxies. Galaxies that possess discs but not 
spiral arms are called lenticular galaxies, because they look lens-shaped when seen 
edge-on. 

The classification system is further subdivided and specialized to take account 
of, for instance, the brightness of the nuclear region (the tight compact central 
region of the galaxy), the tightness of the spiral arms, and so on. 


4.3 Galaxy Structure 


At this point, I think it will be helpful to describe in a little more detail the 
structure of a galaxy, which will, albeit in a small way, provide some insight into 
why galaxies appear the way they do. The books mentioned in the appendix will 
have a much more detailed coverage of this topic along with discussions on the 
origin and formation of galaxies. 

Spiral galaxies have a thin disc extending outward from a central bulge. The 
bulge merges smoothly into the halo, which can extend to a radius in excess of 
100,000 l.y. Both the bulge and halo make up the spheroidal component. There 
are no clear boundaries as to what divides this component up into its constituent 
parts, but a ball-park figure often used is that stars within 10,000 l.y. of the 
center can be considered to be bulge stars, whereas those outside this radius are 
members of the halo. 

The disc component of a spiral galaxy cuts through both the halo and the 
bulge, and can, in a large spiral galaxy such as the Milky Way, extend 50,000 l.y. 
from the center. The disc area of all spirals contains a mixture of gas and 
dust, called the interstellar medium, but the amounts and proportions of the 
gas, whether atomic, ionized, or molecular, will be different from galaxy to 
galaxy. 





Galaxies 


159 


4.4 Stellar Populations 


The stars contained within a spiral galaxy can also be classified by where they 
reside. Those that lie in the disc region are called Population I stars and are often 
young, hot, and blue stars; those in the bulge region are old, red giant stars called 
Population II stars. This is why photographs often show the spiral arms colored 
blue, due to the Population I stars, with the bulge colored orange because of the 
Population II stars. The spiral arms may also be dotted with pink and red HII 
regions, 4 areas of star formation. Thus, new stars are usually formed in the spiral 
arms of galaxies, seldom in the bulge. 

About 75% of large galaxies in the observable universe are apparently spiral 
or lenticular. Some spiral galaxies can be found in a loose collection of other 
spirals — this is known as a group — spread over several million light years. Our 
Galaxy, the Milky Way, is a member of the Local Group. 

Elliptical galaxies differ significantly from spirals in that they do not have 
a significant disc component. Therefore, an elliptical has only the spheroidal 
component. The interstellar medium is also different from that in spirals; it is a 
mixture of low-density, hot, X-ray-emitting gas. Contrary to what you may read 
in some books, ellipticals do possess a little gas and dust, and some have a small 
gaseous disc at their center, which is believed to be the remains of spiral galaxies 
that the elliptical has consumed. 

The stars in the spheroidal population of elliptical galaxies give a clue to possible 
star formation, if any. Such stars are orange and red; the absence of blue stars 
indicates that they are old, and that star formation occurred a long time ago. 

Elliptical galaxies are often found in large clusters of galaxies, usually located 
near their center. They make up about 15% of the large galaxies found outside the 
clusters, but about 50% of the large galaxies within a cluster. Very small galaxies, 
called dwarf elliptical galaxies, are often found accompanying large spiral ones. 
A perfect example of such an arrangement, and one that is visible to an amateur, 
is the Great Andromeda Galaxy, M31, which is a classic spiral galaxy, and its 
attendants, M32 and M110, both dwarf ellipticals. 

Several galaxies can be observed to not belong to either the spiral or the 
elliptical category. These are irregular galaxies, which more or less include all 
those galaxies that do not easily fall into the two previous classes. They include 
small galaxies, such as the Magellanic Clouds, 5 and those galaxies that are peculiar 
due to tidal interactions. These systems of galaxies are usually white and dusty, 
as spirals are, though there the resemblance ends. Deep imaging has shown that 
more distant galaxies are irregular, which indicates that this type of galaxy was 
more common when the universe was much younger. 


4.5 Hubble Glassification 
of Galaxies 


The famous American astronomer Edwin Hubble was the first to put the many 
disparate types of galaxies into some sort of order. The Hubble Classification, 





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as it is now known, was used as a means of categorizing a galaxy. Further 
amendments have, of course, been made, particularly by the astronomer Gerald 
de Vaucouleurs. 

Basically, the classification is as follows. An upper-case letter followed by either 
a number or a lower-case letter is assigned to the galaxy in question, and this 
identifies its morphology. For instance, in the case of an elliptical galaxy, the 
letter E is used followed by a number. The larger the number, the flatter the 
galaxy. An EO galaxy is round, whereas an E7 galaxy is very elongated. There 
exists a subgroup for the ellipticals with the following nomenclature: D signifies 
a diffuse halo, c is a supergiant galaxy, and d represents a dwarf galaxy. Thus, 
some of the largest elliptical galaxies have a cD classification. 

A spiral galaxy is assigned the letter S, but it can also be assigned SA to signify 
that it is an ordinary spiral, or SB where B indicates that it has a bar. It is 
then followed by a lower-case letter: a, b, c, or d. Intermediate classes also exist, 
namely, ab, be, cd, dm, and m. The lower-case letters a to d indicate the size of 
the bulge region, the dustiness of the disc, and the tightness of the spiral arms, 
while m denotes a stage where the spiral shape is barely discernible. An Sa galaxy 
will usually have a large bulge, a modest amount of dust, and tightly wound 
arms, whereas an Sd galaxy will have a small bulge and very loosely wound arms. 
An SBc galaxy will have both a bar and a small bulge. 

Galaxies that are intermediate between spirals and ellipticals, called the 
Lenticular galaxies, are classified as SOs; for instance, SAOs those that are 
ordinary, and SB</>s those that are barred. In addition, for galaxies intermediate 
between types S and SB, there is the classification SAB. Both lenticular and spiral 
galaxies can be surrounded by an outer ring, and the spiral arms can nearly close 
upon themselves, thus forming a pseudo-ring. These new features are classified 
as R and R\ respectively. 

Finally, there are classifications for galaxies that do not easily fall into any of 
the above three! These include Pec, for peculiar galaxies, which have a distorted 
form. Some galaxies have an irregular morphology, classed as Irr. They can 
also be further classified as unstructured, IA, and barred, IB. Dwarf galaxies are 
classified as d. The difference between a Pec and an Irr can be very small, but it 
appears that a peculiar galaxy is one that may have suffered considerable tidal 
distortion for the passage of another galaxy nearby. 

The Hubble classification system can be represented by a simple diagram 
(Figure 4.1); note, however, that the diagram and the classification generally 
do not represent an evolutionary sequence. Galaxies do not start as ellipticals 
and then progress to be spirals, though there is some evidence that the reverse 
is true. 

The classification system can be confusing (an understatement!), and the 
system and descriptions outlined above are by no means complete — there are 
further subdivisions to all the classes. But do not let that worry you — the complete 
system is only of relevance to those astrophysicists who study galaxies; to the 
observer the important point is whether the galaxy is elliptical or spiral and, if 
spiral, it is barred. In a few galaxies where the spiral or elliptical structure is 
very apparent, the subdivisions of, say, El, E2 and Sa, Sb, SBa, and so on, will 
be useful. Like most things in observational astronomy, it will all become easier 
with continued use. 


Galaxies 


161 


Sa Sb Sc 



PorofO 


SBa SBb SBc 

Figure 4.1. Hubble's tuning fork diagram showing the main galaxy types. 


Sometimes the classification system may be of no apparent use at all, such as 
when a galaxy is inclined to our line of sight. The galaxy M83, which is a nice 
spiral galaxy, is face-on to us, and thus the classic spiral shape is very apparent. 
However, the galaxy NGC 891, which is classified as a spiral, will just appear as 
a thin streak of light because the galaxy is edge-on to us, and so the spiral shape 
will not be visible. 

However, finding galaxies that present an edge-on or nearly edge-on 
perspective adds another element to the pleasure of locating and observing these 
faint objects. 


4.6 Observing Galaxies 


To an amateur astronomer, observing galaxies can present something of a 
dilemma. In astronomical magazines and books, you are bombarded with images 
of galaxies, their spiral arms resplendent and multi-colored, speckled throughout 
with distinct pink HII regions. However, when you look at that same galaxy 
through your telescope, all you can see is a pale tiny blob! 

It is true that in nearly every case, especially from an urban location, the 
galaxy will be faint and indistinct. Only with the largest telescopes and the 
darkest possible skies can you see any real structure. But take heart! You will be 
astonished at what you can actually see with the naked eye with practice from 
the right location. I recall that, during my visit to the wilder parts of Turkey, on 
several occasions under utterly dark skies, I was able to see M31 in Andromeda 
and M33 in Triangulum in such amazing detail that even today the memory takes 
my breath away. With just my naked eye, I was able to trace M31 to nearly 2 V 2 0 
across the sky, and M33 was a huge amorphous glow. Also, had I known enough 



Astrophysics is Easy 


162 

about how to look, I would have been able to see several other galaxies with 
the naked eye; unfortunately I was under the common misapprehension that the 
naked-eye limit is about 6th magnitude, but I now know that, with extremely 
dark skies and light-adapted vision, magnitude 8 is more likely the limit. 

The purpose of this anecdote is to remind you that, in order to see faint 
galaxies and detail therein, dark skies are indispensable. With very dark skies, 
and armed with just a pair of binoculars, many galaxies are within reach. If you 
have a telescope, the number increases dramatically. 

As usual, dark skies and dark- adapted and averted vision will help in tracking 
down and seeing galaxies. Clean optics will also greatly aid you in your obser- 
vations. Dust and smears of grease will reduce a surprising amount of light that 
reaches your eyes, and in particular will reduce the contrast. 

Generally, galaxies that have a brightness greater than 13th magnitude are usually 
visible through telescopes of aperture 15 cm, and telescopes of aperture 30 cm 
will help you see down to about 14.5 magnitude. There are, of course, galaxies 
with much brighter magnitudes than these, and so they can be visible with much 
smaller instruments. In some cases, only the brightest part of a galaxy will be 
visible — perhaps its core (nuclear region), and the spiral part unobservable. 

To be able to trace the finer details of the spiral arms of galaxies, and to locate 
the bulge area, faint halo, and HII regions, you will invariably need a large- 
aperture telescope. However, if the purpose of your observing is just to locate 
these elusive objects, and to be amazed that the light that is entering your eye 
may have begun its journey over 100 million years ago, then there is a plethora 
of galaxies awaiting you. 6 

The usual nomenclature applies in the following descriptions, but with some 
changes. Galaxies are extended objects, which means that they cover an appre- 
ciable part of the sky: in some cases a few degrees, in others only a few arc 
minutes. The light from the galaxy is therefore “spread out,” and thus the quoted 
magnitude will be the magnitude of the galaxy as if it were the “size” of a star; this 
magnitude is often termed the integrated magnitude. This can cause confusion, 
as a galaxy with, say, a magnitude of 8 will appear fainter than an 8th-magnitude 
star, and in some cases, where possible, the surface brightness of a galaxy will 
be given. This will give a better idea of what the overall magnitude of the galaxy 
will be. For instance, Messier 64, the Black Eye Galaxy, has a magnitude of 
8.5, whereas its surface brightness is 12.4. The surface brightness will be given 
in italics after the quoted magnitude; for example, the magnitude and surface 
magnitude of M64 will appear like this: 8.5m [12.4m], 

Following on from the previous paragraph, the designation “easy,” “moderate,” 
or “difficult” takes into account not only the brightness of the galaxy but also 
the area of the sky the galaxy spans. Thus, a galaxy may be bright with, say, 
a magnitude of 8, which, under normal circumstances, would be visible with 
binoculars and designated as “easy”; however, if it covers a significant amount 
of the sky (and thus its surface brightness is low, making it more difficult to 
observe) I would designate it “moderate.” 

In addition, spiral galaxies can exhibit a variety of views, depending on their 
inclination to the Solar System. Some will appear face-on, others at a slight 
angle, and a few completely edge-on. As an indicator of inclination, the following 
symbols will be used: 


Galaxies 


163 


Face-on: 

Slight inclination: 

Edge-on: 

Finally, the Hubble classification of galaxies I outlined earlier will also be used. 

4.6.1 Spiral Galaxies 

Caldwell 7 NGC 2403 07 h 36.9 m +65°35' January 14 

8.5m [13.9m] 17.8' | 10.0' b* SAB(s)cd easy 

This is one of the brightest galaxies, which was omitted from the Messier 
catalogue, and is often left out of an observer’s schedule. When observed through 
binoculars, it appears as a large, oval hazy patch with a brighter central region. 
With averted vision and an aperture of about 20 cm, faint hints of a spiral 
arm will become apparent. Larger apertures will, of course, present even further 
detail. It is not a member of the Local Group of Galaxies 7 but believed to be a 
member of the M81-M82 group. It was the first galaxy outside the Local Group 
to have Cepheid 8 variable stars discovered within it, and the current estimate of 
its distance is 11.5 million l.y. 

Messier 81 NGC 3031 09 h 55.6 m +69°04' February 18 

6.9m [13.0m] 26' 1 14' •_ SA(s)ab easy 

A spectacular object! With binoculars, it will show a distinct oval form, and with 
high-power binoculars, the nuclear region will easily stand out from the spiral 
arms. A telescope can show considerably more detail, and it is one of the grandest 
spiral galaxies on view. With an aperture of about 15 cm, traces of several of the 
spiral arms will be glimpsed. A real challenge, however, is to locate this galaxy 
with the naked eye. Several observers have reported seeing it in dark skies. If you 
happen to glimpse it without any optical aid, you are probably looking at one of 
the farthest objects 9 that can be seen with the naked eye, lying at a distance of 
some 4.5 million l.y. M81 is partner galaxy to M82, and both of these spectacular 
objects can be glimpsed in the same field of view. 

Caldwell 48 NGC 2775 09 h 10.3 m +07° 02' February 7 

10.1m [13.1m] 5.0' 1 4.0' SA(R)ab moderate 

A difficult object for binoculars, this galaxy is observable only through a telescope. 
With an aperture of about 20 cm, you will see the galaxy as a large blob. However, 
detail within the object is conspicuously absent, but a brighter core and fainter 
outer region can be resolved. The absence of detail (e.g., spiral arm dust and 
gas) has been attributed to an early era of star formation, which used up all 
the material. The evidence for this was found in the galaxy’s spectrum, which 
that lacked emission lines because these lines are usually caused by star-forming 
regions in and around the spiral arms. 

Messier 96 NGC 3368 10 h 46.8 m +11°49' March 3 

9.2m [12.9m] 7.6' 1 5.2' w? SAB(rs)ab easy 


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Astrophysics is Easy 


This faint galaxy can be seen through binoculars as a faint, hazy oval patch of 
light. What you would observe is, in fact, just the bright central nucleus of the 
galaxy, as the spiral arms are too faint to be resolved. Telescopes will bring out 
further detail, and with good conditions the spiral arm features can be seen. 
There is some controversy over M96, as recent measurements of its distance 
show that it is 38 million l.y, which is 60% greater than the previous value. It 
forms a nice triangle with two other galaxies, M95 and M105. 

Messier 65 NGC 3623 ll h 18.9 m + 13°05' March 12 

9.3m [12.4m] 9.8' 1 2.9' SAB(rs)a easy 

Visible through binoculars, Messier 65 is one half of the most famous galaxy 
pair in the sky, after M81 and M82. Along with M66, it shows up quite well 
with low-power optics. It appears as a nice oval patch of light, and with higher 
magnification both spiral arms and a dust lane can be glimpsed. However, it is 
often difficult to observe, as the brightness of its background tends to interfere 
with the galaxy details. A nice challenge for observers in an urban location. 

Messier 66 NGC 3627 ll h 20.2 m +12°59' March 12 

9.0m [12.5m] 9.1' |4.2' SA(s)b easy 

The other half of the galaxy duo mentioned above. This is a bright galaxy, easily 
seen through binoculars, where its distinct elliptical shape and bright center can 
be resolved. With telescopes, the oval shape of the nucleus becomes apparent, 
and with higher magnification a spiral arm and dark patch can be seen. Large- 
aperture telescopes will show considerable detail, consisting of dark and light 
patches. 

Messier 106 NGC 4258 12 h 19.0 m +47°18' March 27 

8.3m [13.8m] 18.6' 1 7.2' SAB(s)bc easy 

The galaxy appears as a large glow when seen through binoculars and has a 
distinct elliptical shape. Large binoculars reveal the presence of the nucleus. 
With telescopes of small aperture (10 cm) and low magnification, the spiral arms 
become apparent, and with higher magnification further detail can be seen. This 
galaxy needs a large aperture and magnification to reveal any amount of detail. 
The galaxy is nearly face-on to us, but a cloud of gas and dust surrounding the 
nucleus is apparently edge-on. Furthermore, there is evidence that a black hole 
resides at the core! 

Messier 88 NGC 4501 12 h 32.0 m +14°25' March 30 

9.6m [12.6m] 6.9' 1 3.7' SA(rs)b easy 

It is a fine galaxy, observable through binoculars, with a bright nucleus and a 
faint hazy glow surrounding it caused by the spiral arms. However, a problem is 
that the galaxy is located in a barren patch of the sky, making its location difficult. 
But, with a telescope of medium aperture, say 20 cm, its structure becomes visible. 
The spiral arms and nucleus can be resolved with averted vision. 

Caldwell 38 NGC 4565 12 h 36.3 m +25°59' March 31 

9.6m [13.5m] 15.5' 1 1.9' SA(s)bsp easy 


Galaxies 


165 


It is a striking example of an edge-on galaxy; with small binoculars, the classic 
spindle shape can be seen against the background of stars, and with large ones, 
the central core region can be seen. With a telescope of aperture 15 cm, the lovely 
edge-on shape becomes even clearer, along with its star-like nucleus. The dust 
lane is observable only with apertures of at least 20 cm. It is thought to be a 
massive galaxy, similar to the Milky Way, with its dust lane the equivalent of 
The Great Rift. 

Messier 98 NGC 4192 12 h 13.8 m +14°54' March 25 

10.1m [13.2m] 9.8'|2.8' SAB(s)ab moderate 

This faint galaxy lies at the edge of the great Coma-Virgo Cluster, an area studded 
with galaxies both faint and bright. It is a difficult object to locate and requires 
a small telescope of at least 10 cm aperture. It is highly inclined to us and has 
a very elongated shape. With excellent seeing conditions, a higher magnification 
would show spectacular spiral arms and dust lanes, and rarely the entire halo can 
be seen to surround the galaxy. It is an excellent object to observe but requires 
skill and patience to see any detail in it, especially from an urban location. 

Messier 99 NGC 4254 12 h 18.8 m +14°25' March 27 

9.9m [13.0m] 5.4' 1 4.7' SA(s)c moderate 

A difficult object for binoculars; its circular disc shape can be resolved only in 
dark skies. With a telescope of aperture 10 cm, the galaxy will remain a hazy 
round patch, but a small nucleus may be visible. Higher magnification and 
perhaps greater aperture will show two spiral arms. However, M99 is one of 
those galaxies whose likelihood of showing detail depends greatly on the seeing 
conditions. Try observing on a clear night and then on an average night, and 
compare your observations. M99 lies at a distance of about 55 million l.y. and is 
one of the galaxies within the Coma-Virgo Cluster. 10 

Messier 61 NGC 4303 12 h 21.9 m +04°28' March 28 

9.6m [13.4m] 6.5'|5.8' SAB(rs)bc moderate 

A difficult object for binoculars, you would see nothing more than a small faint 
circular patch of light even in dark skies. For small-telescope users, however, it 
is a delight, although it is small and difficult to locate. It is an ideal open-faced 
spiral galaxy. The use of averted vision is a must for this object, when the nucleus 
and any spiral arm detail become apparent. A nice addition is the fact that the 
galaxy is located within the Virgo Cluster of galaxies, and you may notice several 
faint and indistinct glows in the same field of view as M61, depending on the 
limit of your vision. These are probably unresolved galaxies! 

Caldwell 36 NGC 4559 12 h 35.9 m +27°57' March 31 

9.8m [13.9m] 13.0' 1 5.2' " SAB(rs)cd moderate 

A member of the Virgo Cluster, this is often overlooked. Not really a binocular 
object, it is perfect for small telescopes. With a 20 cm aperture, the clear oval 


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shape will be resolved, along with a brightening of the nucleus. A perceptible 
hint of further detail may also be glimpsed. Larger apertures show considerably 
more detail. 

Caldwell 40 NGC 3626 ll h 20.1 m +18°21' March 12 

11.0m [12.8m] 2.8' 1 2.0' (R)SA(rs) difficult 

This galaxy is virtually unknown among amateurs. It lies close to several brighter 
Messier objects, and is often mistaken for NGC 3607. Nevertheless, it is worth 
searching for. It is a featureless oval patch of light, and you will need an aperture 
of at least 20 cm to resolve this object. What makes it special, however, is that 
it is a multispin galaxy. This means the molecular and ionized gas is rotating 
around the galaxy in the direction opposite to that of its stars. The origin of this 
phenomenon is unknown, but one school of thought suggests that the galaxy 
recently collided, assimilating a huge gas cloud with a mass of about one billion 
solar masses. 

Caldwell 26 NGC 4244 12 h 17.5 m +37° 48' March 26 

10.4m [14.0m] 18.5' |2.3' SA(s)cd:sp difficult 

This is one of those needle-like galaxies that can be seen with amateur instru- 
ments. The edge-on galaxy is exceedingly thin, and you will need an aperture 
of at least 20 cm to locate and observe it. It has a faint but easily resolvable 
star-like nucleus, but detail within the galaxy is very rarely seen, even with larger 
apertures. It has a Hubble classification similar to that of M33, the Pinwheel 
Galaxy in Triangulum. When such a galaxy is seen edge-on, the tiny nucleus and 
loose, open arms give it its indistinct appearance. 

Messier 58 NGC 4579 12 h 37.7 m +11°49' April 1 

9.6m [13.0m] 5.9' 1 4.7' SAB(rs)b easy 

This galaxy will appear as a faint, hazy patch of light with a barely discernible 
nucleus when observed through binoculars. You may also glimpse in the same 
field of view the galaxies M59 and M60. A telescope with aperture 10 cm will show 
some structure in the halo, along with faint patches of light and dark. There are 
some reports that a 20 cm telescope will allow the bar connecting the spiral arms 
to the nucleus to be resolved. It has about the same mass as the Milky Way and 
is about 95,000 l.y. in diameter. 

Messier 104 NGC 4594 12 h 40.0 m -11°37' April 1 

8.0m [11.6m] 8.7' 1 3.5' SA(s)asp easy 

Also known as the Sombrero Galaxy, it is an extragalactic treasure! It gives a 
marvellous sight with almost all binoculars and telescopes. With small binoculars, 
it will reveal itself as an oval disc, which increases in brightness toward the center. 
Large binoculars will, however, reveal its true beauty. The dark dust lane that 
cuts across the galaxy becomes readily apparent. Looking though a telescope, 
even more detail can be brought out. With a 10 cm aperture, a bright core can 


Galaxies 


be seen, along with the long, spindle-like dust lane. With higher magnification, 
the spiral arms stand out. Large apertures and higher magnifications will reveal 
a wealth of detail. It was the first galaxy, other than the Milky Way, to have its 
rotation determined. 

Messier 94 NGC 4736 12 h 50.9 m +41°07' April 4 

8.2m [13.0m] 11.2'|9.1' ^ •- — (R)SA(r)ab easy 

This galaxy is visible through binoculars and will appear as a small circular hazy 
patch with a star-like nucleus. Telescopes will reveal some structure, possibly 
a faint spiral arm. Several observers have reported that a central ring can be 
seen near the core, which gives it an appearance very similar to M64, the Black 
Eye Galaxy. A further elliptical ring has been reported outside the edge of the 
galaxy. Needless to say, exceptionally dark skies will help one observe this elusive 
feature. 


Messier 64 NGC 4826 12 h 56.7 m +21°41' April 5 

8.5m [12.4m] 9.3'|5.4' (R)SA(rs)ab easy 

Also known as the Black Eye Galaxy, this famous galaxy can be seen through 
small binoculars as an oval hazy patch with a slightly brighter center. The striking 
feature that gave the galaxy its name has been reported to be visible through 
large binoculars on very dark nights. Small telescopes show a very bright nucleus 
encased in a patch of glowing light. The “eye” is a vast dust lane, some 40,000 l.y. 
in diameter. There is considerable debate as to whether the “eye” can be seen 
through small instruments. Some observers report that an aperture as small 
as 6 cm will resolve it, while others claim at least 20 cm is needed. What was 
concluded is that a high-magnification instrument is necessary. There is further 
controversy as to whether the nucleus is star-like or not. It may be that magni- 
fication plays an important role here. 

Messier 63 NGC 5055 13 h 15.8 m +42°02' April 10 

8.6m [13.6m] 12.6' 1 7.2' SA(rs)bc easy 

Also known as the Sunflower Galaxy, this object is somewhat difficult for binoc- 
ulars, even though it has a fairly bright magnitude. In small binoculars, it will 
just appear as a faint patch of light, and with large binoculars, the classic oval 
shape will become apparent. Telescopes reveal a lot of detail, including many 
faint and detailed spiral arms. 

Messier 51 NGC 5194 13 h 29.9 m +47° 12' April 14 

8.4m [13.1m] 11.2'|6.9' SA(s)bcP easy 

Also known as the Whirlpool Galaxy, this famous galaxy is easily visible with 
binoculars. It appears as a small glowing patch of light with a bright, star-like 
nucleus. Many now believe that M51 is the finest example of a face-on spiral 
galaxy. What makes this galaxy so special is the small, irregular, satellite galaxy 
NGC 5195, which is close to it. Deep photographs reveal that the galaxies are 
physically connected with a bridge of material, but unfortunately this satellite 


168 


Astrophysics is Easy 


galaxy cannot be seen through most binoculars; even with giant binoculars, 
it appears only as a slight bump on the side of M51. With small telescopes 
(10 cm), not a lot of detail is visible, except perhaps the slightest hint of spiral 
structure. With an aperture of 25 cm, much more detail can be resolved: spiral 
arms, structure within the arms, and dark patches. A matter of debate, however, 
is whether the bridge connecting M51 to NGC 5195 can be seen with small 
telescopes. Some observers claim that it can be seen with 10 cm aperture; others 
claim that at least 30 cm is needed. What everyone would agree upon is that 
absolute, perfect transparency is needed, as even the slightest haze or dust in the 
atmosphere will make observations much more difficult. 

Messier 83 NGC 5236 13 h 37.0 m -29°52' April 16 

7.5m [13.2m] 12.9'| 11.5' SAB(s)c easy 

Often overlooked by observers, this is a nice galaxy located within the star fields 
of Hydra, and it is a showpiece for small telescopes. With binoculars, it will 
appear as a hazy patch of light with a bright, star-like nucleus. With telescopes, 
much more detail can be seen, including spiral arms, dust lanes, bright knots, 
and even detail within the nucleus. It is one of those objects that merits several 
observing sessions. Also, it ties with M81 as one of the farthest objects visible to 
the naked eye, at about 4.5 million l.y. As it is located too far to the south, there 
may be a problem for observers from the UK in locating it. 

Caldwell 30 NGC 7331 22 h 37.1 m +34°25' August 30 

9.5m [13.5m] 1 1.4' |4.0' SA(s)bc easy 

This is the brightest galaxy in Pegasus. With binoculars, it will appear as a faint 
patch of light with a brighter core. A telescope of 20 cm will show its structure 
in a little more detail. Apparently, this galaxy is similar to M31 but lies much 
farther from us, at a distance of 50 million l.y. There is also some debate as to 
whether this galaxy is linked with the famous Stephen’s Quintet (see the section 
on “Groups and Clusters of Galaxies”). 

Caldwell 43 NGC 7814 00 h 03.2 m +16°08' September 21 

10.6m [13.3m] 6.3' 1 3.0' SA(s)ab:sp moderate 

This is no binocular object, but nevertheless a splendid sight with larger-aperture 
telescopes. It is a fine example of an edge-on galaxy and bears many similarities 
to its better-known cousin, Ml 04. It can be easily seen with a telescope of aperture 
20 cm, but it often provokes debate among amateurs as to whether its dust lane 
can be seen with small telescopes. Some profess to have seen it with 20 cm, but 
others claim that at least 40 cm is needed. Try observing with as high a power it 
can take, as this may help you resolve this dilemma. 

Messier 31 NGC 224 00 h 42.7 m +41° 16' October 1 

3.4m [13.6m] 3°| 1° Sb easy 


Galaxies 


169 


Also known as the Andromeda Galaxy, this most famous galaxy in the sky is also 
the most often visited one and is always a first observing object for the beginner. 
It can be visible to the naked eye, even on those nights when the conditions 
are far from perfect. Many naked-eye observers claim to have seen the galaxy 
spread over at least 2Vi° of sky, but this depends on the transparency. With 
binoculars, it presents a splendid view, and the galactic halo is easily seen, along 
with the bright nucleus. Large binoculars may even show one or two dust lanes. 
Using averted vision and a very dark sky, several amateurs report that the galaxy 
can be traced to about 3° of the sky with telescopes of aperture 10 cm. With 
larger telescopes, a wealth of detail becomes visible. With an aperture of about 
20 cm, a star-like nucleus is apparent, cocooned within several elliptical haloes. 
Another striking feature is the dust lanes, especially the one running along its 
north-western edge. Many observers are often disappointed when they observe 
M31, as the photographs they see in books actually belie what they can see at 
the eyepiece. M31 is so big that any telescope cannot really encompass all there 
is to show. Observing this wonderful galaxy with patience will reward you with 
a lot of surprises. Spend several nights observing the galaxy, and choose a dark 
night in a country location. It really is a spectacular galaxy. It contains about 300 
million stars with a diameter of 130,000 l.y., and it is among the largest galaxies 
known. It is the largest member of the Local Group. In older texts, it is often 
referred to as the Great Nebula in Andromeda. 

Caldwell 65 NGC 253 00 h 47.6 m -25° 17' October 3 

7.2m [12.6m] 25.0' 1 7.0' SAB(s)c easy 

This is a wonderful object and is often referred to as the southern hemisphere’s 
answer to the Andromeda Galaxy. It can be easily seen with binoculars as a long, 
spindle-shaped glow with a bright nucleus. With large binoculars, some structure 
can be glimpsed under absolutely perfect conditions. Its size makes it impressive, 
as it is as wide as the Moon and around a third as thick. Any telescope, even 
as small as 6 cm, will suffice to see this object. Larger apertures will reveal more 
detail, and with averted vision the spiral arms can be glimpsed. Several reports 
suggest that considerable mottling can be seen within the galaxy with a 15 cm 
telescope. It has the dubious honor of being one of the dustiest galaxies known, 
as well as one that is undergoing a period of frenetic star formation in its nuclear 
regions. 


Caldwell 70 NGC 300 00 h 54.9 m -37°40' October 4 

8.1m [14.7m] 20.0' 1 15.0' C:? SA(s)d easy 

A difficult object to locate due to its very low surface brightness, it will present 
a considerable challenge with binoculars. Nevertheless, once located, it can be 
an impressive object. With a 20 cm-aperture telescope, the nucleus can readily 
be seen engulfed in the unresolved haze of the spiral arms. Larger telescopes 
will resolve some further detail. This galaxy lies at a close distance of about 
7 million l.y. 

Messier 33 NGC 598 01 h 33.9 m +30°39' October 14 

5.7m [14.2m] 71'|42' £■>* SA(s)cd easy 


170 


Astrophysics is Easy 


Also known as the Pinwheel Galaxy, it is a famous galaxy for several reasons. 
It is, without doubt, one of the most impressive examples of a face-on spiral. 
However, it is also one of the most difficult galaxies to locate. Many amateur 
astronomers have never seen it, but others have had no trouble locating it. The 
problem arises from its having such a large surface area. Although it has an 
integrated magnitude of 5.7, it spreads the light out to such an extent that it 
appears very faint. As a result, the galaxy may be all but invisible with telescopes, 
whereas it can be easily seen with binoculars. It will look like a large, very faint 
cloud with a slight brightening at its center. In addition, there are several reports 
of its being visible to the naked eye; I testify to this, and it was strikingly visible 
from a totally dark night under perfect conditions, when it was impossible to 
see it otherwise! With a telescope of aperture 10 cm, several spiral arms can be 
glimpsed arcing from the very small nucleus. With large telescopes, a plethora 
of detail becomes visible, such as star clusters, stellar associations, and nebulae, 
all located within the galaxy. 11 This truly is a spectacular galaxy. 

Caldwell 62 NGC 247 00 h 47.1 m -20°45' October 2 

9.1m [14.1m] 20.0'|7.0' ^ SAB(s) moderate 

This galaxy is barely discernible with large binoculars, where it is seen as an 
elongated hazy patch of light with a brighter nucleus. However, it lies low down 
in the skies, and so is often neglected by UK observers. With larger-aperture 
telescopes, its mottled appearance becomes visible, along with the brighter, 
southern part of the galaxy. The northern part is much fainter and will require 
averted vision and clear skies for observation. The galaxy was thought to be a 
member of the Sculptor Group of galaxies, but recently doubts have arisen, as 
the most recent estimates of its distance put it at about 13.5 million l.y., which 
is twice the distance of the cluster. 

Caldwell 23 NGC 891 02 h 22.6 m -42°20' October 27 

9.9m [13.8m] 14.0' 1 3.0' SA(s)sp moderate 

This is a fine example of an edge-on galaxy, and it is thought by many to be the 
finest galaxy. With binoculars, it is just visible as a hazy but distinct elongated 
smudge. With a telescope of aperture 20 cm, its spindle shape is very apparent, 
and with a larger aperture the distinctive dust lane will be resolved. 


4.6.2 Barred Spiral Galaxies 

Messier 95 NGC 3351 10 h 44.0 m +11°42' March 3 

9.7m [13.5m] 7.4' 1 5.0' w? SB(R)b moderate 

This is a faint galaxy, which shows little, if any, detail in binoculars. It will just 
appear as a hazy patch, lying in the same field of view as M96. With a telescope 
of at least 15 cm, some structure can be glimpsed, with larger apertures showing 
the distinctive bar feature. There is some debate as to the real magnitude of the 
galaxy, with some observers putting it at 9.2 m. 


Galaxies 


171 


Messier 108 NGC 3556 ll h 11.5 m +55° 40' March 10 

10.0m [13.0m] 8.7' 1 2.2' SB(s)cdsp moderate 

This galaxy is visible with binoculars as a very faint streak of light. The central 
condensation has been reported to be visible with an 8 cm telescope. Larger 
apertures show a surprising amount of detail with considerable mottling and 
structure. This is a very small galaxy, just one-twentieth the mass of the Milky Way, 
lacking a central bulge. Although it was a recent addition to Messier’s list, he was 
aware that the galaxy existed, but for some reason, he just did not include it. 

Messier 109 NGC 3992 ll h 57.6 m +53°23' March 21 

9.8m [13.5m] 7.6'|4.7' ■«il— ■ SB(rs)bc moderate 

Another recent addition to the Messier catalogue, this galaxy can be glimpsed 
with binoculars, providing the conditions are right. With low-power and small- 
aperture telescopes, it is evident that you are looking at a galaxy, but no further 
detail can be seen. High-power and larger-aperture telescopes can show some 
structure, such as the core and halo regions. The central bar also becomes 
prominent with aperture around 25 cm. This is the penultimate Messier object. 

Caldwell 3 NGC 4236 12 h 16.7 m +69°27' March 26 

9.6m [14.7m] 23.0' |8.0' SB(s)dm difficult 

Although this is a large galaxy, it is very faint and so difficult to locate. In 
addition, as it is edge-on to us, it presents a very slim view, and so spiral arm 
features are absent. With aperture around 20 cm, its distinctive spindle shape is 
conspicuous. It is a nice galaxy for those who would like to test the limits of 
a small telescope as well as their observing skill. It lies at a distance of about 
10 million l.y. 

Messier 91 NGC 4548 12 h 35.4 m +14°30' March 31 

10.1m [13.3m] 5.4' 1 4.3' SB(rs)b difficult 

Now it is time for a mystery! If you try to locate M91 from Messier’s original 
notes, you may make an interesting discovery. There is nothing there! Most 
observers agree that Messier made a mistake, and that in fact the galaxy NGC 
4548, what he originally observed, was incorrectly plotted. The galaxy is a faint 
object, and telescopes of medium aperture will be needed to see any detail; yet 
it may be visible with aperture of 10 cm as a faint, hazy circular patch. 

Caldwell 32 NGC 4631 12 h 42.1 m +32°32' April 2 

9.2m [13.3m] 17.0' 1 3.5' SBfsjdsp moderate 

Surprisingly an often neglected galaxy, it has a lot to offer. Visible with binoculars 
as a faint elongated object, you will really need a telescope to appreciate its 
beauty. It is a very big galaxy; due to its appearance, it was unofficially nicknamed 
the Whale Galaxy. Its eastern end is appreciably thicker than its western end, 


Astrophysics is Easy 


172 

hence the name. This aspect can be seen with an aperture of 20 cm, and larger 
telescopes will show further detail, such as patches of light and dark, along 
with two prominent knots. On the northern side of the galaxy is a faint 12th- 
magnitude star, which, providing the seeing is good, will act as a pointer to a 
faint companion galaxy. Several theories have arisen as to the cause of its strange 
and disturbed appearance. The most probable reason is tidal interactions with 
several nearby galaxies. 

Caldwell 72 NGC 55 00 h 15.1 m -39° 13' September 24 

7.9m [13.5m] 25.0'|4.1' SB(s)m:sp easy 

Although this galaxy lies so far to the south it is invisible from the northern 
hemisphere, it still warrants inclusion. With binoculars, it appears as a faint 
spindle-shaped object, and large binoculars hint at some delicate structure. 
Telescopes show even more detail, and it is one of the few galaxies for which an 
H-alpha filter will highlight its HII regions. 

Caldwell 44 NGC 7479 23 h 04.9 m +12° 19' September 7 

10.9m [13.6m] 4.4'|3.4' SB(s)c moderate 

This faint galaxy can be glimpsed with a small telescope of aperture 8 cm as a 
smudge, but do not expect any more detail. With aperture around 20 cm, the 
central bar will be visible, along with a suggestion of some structure. The spiral 
arms at the end of the bar will need at least a 30 cm telescope, but some observers 
claim to have seen them with a 25 cm aperture under perfect seeing conditions. 
The nucleus can be easily resolved, however. It has the honor, among some 
amateur astronomers, of being the finest barred spiral on offer for the northern 
hemisphere. It lies at about 100 million l.y. from us. 

NGC 1365 03 h 33.6 m -36° 08' November 14 

9.5m [13.7m] 9.8' 1 5.5' (R)SB(s)b moderate 

This is a very impressive galaxy, easily visible with binoculars as an elongated 
hazy object with a brighter center. With a telescope having aperture as small 
as 8 cm, its origin is obvious, and larger apertures will show considerably more 
detail. Although not visible from the UK, it should be a nice observing target 
from the USA. 

4.6.3 Elliptical Galaxies 

Messier 84 NGC 4374 12 h 25.1 m +12°53' March 28 

9.1m [72.3m] 6.5' 1 5.6' El easy 

Located close to the Virgo Cluster of galaxies, it presents a small oval patch of 
light when seen through binoculars. The bright nuclei can be glimpsed under 
favorable conditions. As with most ellipticals observed with amateur telescopes, 
there is never much detail seen in the galaxy; most remain as smooth objects 
with little structure and perhaps only a brightening of the core is all that is ever 


Galaxies 


173 


resolved. Nevertheless, they make observing objects worthwhile. There seems to 
be some debate as to whether M84 is in fact an El galaxy or an SO, which is a 
galaxy between an elliptical and a spiral. It is located in the Virgo Cluster at about 
55 million l.y. The area around M86 is full of very faint galaxies, and although 
only a handful will show any perceptible detail, it is nevertheless worthwhile 
sweeping the area for these most elusive objects. 

Messier 86 NGC 4406 12 h 26.6 m +12°57' March 29 

8.9m [13.9m] 8.9' 1 5.8' E3 easy 

A companion and very similar in appearance to M84, this is the brighter of the 
two, with M86 being perhaps slightly brighter with a less condensed core. It is 
generally visible with binoculars and telescopes of all sizes. Similar to the above, 
the area is full of galaxies, and with patience, many of them can be seen in a 
dark sky. 


Messier 49 NGC 4472 12 h 29.8 m +08°00' March 30 

8.4m [12.9m] 10.2' 1 8.3' E2 easy 

This is the second-brightest galaxy in Virgo. It can be easily spotted with 
binoculars as a featureless, oval patch of light. Although most ellipticals are 
rather featureless, M49 stands out quite well with higher magnification and large 
aperture, and some resolution can be seen in the nuclear area too. It seems 
to have a bright nucleus surrounded by a diffuse core region, which in turn 
is surrounded by a rather diffuse halo. Some observers report that the nucleus 
shows a mottled appearance under magnification. The galaxy is at the center of 
a subcluster of galaxies called the Virgo Cloud, which in turn is part of the much 
larger Virgo Cluster. In addition, the elliptical galaxy is cocooned in an envelope 
of hot gas at a temperature of about 10,000,000 K. At such a high temperature, 
X-rays are formed, which may be detected with an X-ray telescope. 

Messier 89 NGC 4552 12 h 35.7 m + 12°33' March 31 

9.7m [12.3m] 5.1' |4.7' E0 moderate 

It is a difficult galaxy to locate with binoculars, especially if the seeing conditions 
are far from ideal. If, however, located, it would appear as just a small hazy spot 
of light. With a telescope and medium magnification, a bright and well-defined 
nucleus can be seen enveloped by the mistiness of the halo. With a large aperture 
and magnification, some mottling has been reported on the halo, but again, the 
atmospheric conditions may limit observability. In the same field of view as M89 
is the spiral galaxy M90. Both are members of the Virgo Cluster. 

Messier 59 NGC 4621 12 h 42.0 m +11°39' April 2 

9.6m [12.5m] 5.4' 1 3.7' E5 moderate 

Although visible with binoculars, this galaxy will pose a challenge to most 
observers. It will probably need the use of averted vision to be spied, and dark 
adaption will undoubtedly be needed. However, telescopically, M59 is rather nice, 
as it is one of the few ellipticals that seem to show detail. It has a star-like nucleus, 


174 


Astrophysics is Easy 


and some observers report a faint mottled appearance, although it could be an 
effect of foreground stars being seen against the oval of the galaxy. It would be 
interesting to find out whether this is correct. Try observing it under excellent 
conditions to see if you can detect any features. Also in the same field of view is 
the elliptical galaxy M60. 

Caldwell 52 NGC 4697 12 h 48.6 m -05°48' April 3 

9.2m [12.7m] 6.0'|3.8' E6 moderate 

This is a nice galaxy, but often ignored and left out from most observing 
schedules. Although it is rather bland in appearance, it stands out well against 
the background star field. Not really a binocular object, it is featureless telescop- 
ically, but with a large aperture, some brightness can be seen at its core. It is a 
dominant member of a small cluster of galaxies, which lies at a distance of about 
60 million l.y. 

Caldwell 35 NGC 4889 13 h 00.1 m +27°58' April 6 

11.5m [13.4m] 2.8'|2.0' E4 difficult 

This is worth seeking out, as it is at a distance of about 350 million l.y. With 
a telescope of at least 20 cm aperture and excellent seeing conditions, you can 
glimpse this tiny object. It has a bright core, surrounded by the usual faint halo. 
It is a dominant member of the Coma Galaxy Cluster, which contains about 
1000 galaxies (several of which can be seen through large-aperture telescopes 
of at least 40 cm). The cluster itself is made of many elliptical galaxies and 
SO-type galaxies. Apparently, this galaxy was the result of a merger of two 
older clusters. Observing any of these galaxies is a feat indeed, but worth the 
effort. 


Caldwell 18 NGC 185 00 h 38.9 m +48°20' September 30 

9.2m [14.3m] 12'|10' dEO moderate 

This is another companion galaxy to M31, as mentioned earlier. However, this is 
easier to locate and observe than M31. With a telescope of 10 cm, it can just be 
glimpsed, whereas with 20 cm, it is easily seen. It remains featureless even with 
larger apertures (40 cm), but shows a perceptibly brighter core. Several reports 
suggest that, with a very large aperture of 75 cm, some resolution of the galaxy 
becomes apparent. It is a dwarf elliptical galaxy. 

Caldwell 17 NGC 147 00 h 33.1 m +48°30' September 29 

9.5m [14.5m] 13'|8.1' dE4 difficult 

Located in Cassiopeia, this is classified as a dwarf elliptical galaxy. Although not 
for from M31, the Andromeda Galaxy, it is in fact a companion to that famous 
galaxy. It is difficult to locate and observe except in very dark skies. It has 
been suggested that a minimum of 20 cm aperture is needed to see this galaxy; 
however, under excellent seeing conditions, a 10 cm telescope is sufficient and 
averted vision is necessary. The moral of this story is that dark skies are essential 
to see faint objects. Increased aperture as well as higher magnification will help 


Galaxies 


175 


its nuclear region become visible. A member of the Local Group, it is one of 
more than 30 galaxies that are believed to be companions to either M31 or the 
Milky Way. 

Messier 110 NGC 205 00 h 40.4 ra +41°41' October 1 

8.0m [13.9m] 21.9'|11.0' E5P easy 

M110 is the final entry in the Messier catalogue, added to the original list in 
1967. It is the second satellite galaxy of M31, and although it has a brighter 
magnitude than that of M32 (the first satellite), it has a much lower surface 
brightness. Consequently, it is much harder to see. However, it is visible with 
large binoculars, but will appear only as a very faint, featureless glow northwest 
of M31. With a telescope, it shows a surprising amount of detail, and a higher 
magnification will bring out its mottled nucleus. In addition, it shows details 
that are peculiar for an elliptical galaxy and visible to the amateur; look for 
dark patches near a bright center. Strangely enough, they are reminiscent 
of features normally found in a spiral galaxy. Of course, exceptionally dark 
skies and perfect seeing and transparency will be needed, but with a telescope 
of even modest aperture, say 10 cm, and a high magnification, they can be 
readily seen. 


4.6.4 Lenticular Galaxies 

Caldwell 53 NGC 3115 10 h 05.2 m -07°43' February 21 

8.9m [12.6m] 8.3' 1 3.2' SOsp easy 

Also known as the Spindle Galaxy, this galaxy is often overlooked, which is 
unfortunate because it is a fine example of its type, as well as being quite bright. 
With binoculars, it will appear as a small, faint elongated cloud, and with large 
binoculars, it will display the characteristic lens shape. It can be easily located 
with telescopes because of its high surface brightness. With telescopes of aperture 
20 cm, it will appear as a featureless oval cloud, with perhaps a slight brightening 
toward its center. As it is classed as an SO-type galaxy, it will not show any 
further detail, even with larger apertures. It is a very big galaxy, about five times 
larger than the Milky Way. It is also one of the most favored objects that is 
purported to have a black hole at its center. 

Caldwell 60/61 NGC 4038/9 12 h 01.6 m -18°51' March 22 

10.3/10.6 m [14.4m] 7.6'|4.9' Sp S(B)p moderate 

They are known as the Antennae or Ring-Tail Galaxies. Together, they probably 
make one of the most famous objects in the entire sky, but few amateurs ever 
observe it, believing it to be too faint. A telescopic object, it will appear as an 
asymmetrical blur through apertures of about 20 cm. Larger apertures will hint 
at its detailed structure, and with a 25 cm aperture, it will resemble the shape 
of an apostrophe. With an aperture of about 30 cm, along with medium or high 
magnification, you will resolve both the objects involved, and it would be a 
worthwhile project to try with a different group of telescopes and observers to 


176 


Astrophysics is Easy 


see just how much detail can be resolved. It is one of those celestial objects that 
are so familiar from photographs that your perception of what you actually see 
will be tainted by what you expect to see. Nevertheless, it is a wonderful object. 
Sadly, it is very low down for UK observers, so perfect observing conditions 
will be necessary. The marvellous shape of the Antennae was the result of spiral 
galaxies passing close by each other, and that tidal interaction caused material to 
be dispersed. Observe the amazing long tails that can be seen from deep images 
of these galaxies. Furthermore, recent work has shown that the interaction has 
led to a vast bout of star formation. Also, it has encouraged astronomers to put 
forward the idea that spiral galaxies evolve into elliptical galaxies after such an 
encounter. 

Messier 85 NGC 4382 12 h 25.4 m +18°11' March 28 

9.1m [13.0m] 7.1'|5.5' SA(s)OP moderate 

This is a bright galaxy that can be glimpsed with binoculars on clear nights. 
With large binoculars, it is even easier, where it will show a star-like nucleus 
surrounded by the faint glow of the halo. A telescope will just magnify the rather 
featureless aspect, though a few observers report that, at high magnification, 
some faint detail can be glimpsed to the south of the nucleus, which may be a 
trace of some spiral structure. Also, there is an indication of a faint blue tint to 
the galaxy. 

Caldwell 21 NGC 4449 12 h 28.1 m +44°05' March 29 

9.6m [12.5m] 6.0'|4.5' IBm moderate 

A member of the Canes Venaticorum Group of galaxies, this is a faint and 

frequently ignored object. Its irregular shape is often mistaken for a comet. 

Under good skies, a telescope of 20 cm aperture will easily discern its fan-shaped 
morphology along with its faint nucleus. Larger telescopes will, of course, resolve 
the galaxy with a considerable amount of detail. An interesting point is that 
several HII regions are visible, especially one at the northern corner of the open 
fan shape. It is apparently a site of much ongoing star formation, and it is similar 
in many ways to the Large Magellanic Cloud. 

Caldwell 57 NGC 6822 19 h 44.9 m -14° 48' July 18 

8.8m [14.2m] 20.0' 1 10.0' IB(s)m moderate 

Also known as Barnard’s Galaxy, this is a challenge for binoculars. Even though 
it is a fairly bright galaxy, it has a low surface brightness and hence is difficult 
to locate. Once located, however, it will just appear as indistinct glow running 
east to west. This is in fact the bar of the galaxy. Strangely enough, it is one of 
those objects that are easier to locate using a small aperture (say 10 cm) rather 
than large. Nevertheless, dark skies are essential to locate this galaxy. 

Caldwell 51 IC 1613 01 h 04.8 m +02°07' October 7 

9.2m [— ] 1 1.0' 1 9.0' dIA moderate 

A very difficult galaxy to observe, a few observers report that it is visible through 
large binoculars as a very faint hazy glow, while others claim that a minimum of 


Galaxies 


177 


20 cm aperture is needed. Whatever you choose, one thing is paramount — a dark 
sky. A member of the Local Group, it is similar in many respects to Caldwell 57. 
It is also an old galaxy that is still forming stars. 


4.7 Active Galaxies and AGNs 


We now discuss a type of galaxy that has become very prominent in astrophysical 
research over the past 20 years or so — Active Galaxies and Active Galactic Nuclei, 
or AGNs. 

The story begins in the 1950s, when radio astronomers began to detect galaxies 
that emitted vast amounts of radio energy, in some cases as much as 10 million 
times more radio energy than a normal galaxy would. Later on, when space- 
borne telescopes made an appearance, it was found that there exist galaxies that 
emit incredible amounts of energy in the infrared, ultraviolet, and X-ray. These 
are the active galaxies. Subsequent observations have shown that most of this 
“extra” energy is emitted from the central regions of the galaxies, and so this is 
the explanation for the term active galactic nuclei. 

The different types of active galaxies are legion! In fact, there was a time when 
it seemed that every time an active galaxy was discovered, it could be put into 
its own individual class. But we now know different. 12 For instance, there are: 13 

Seyfert Galaxies, types 1 and 2 [and 1.6, 1.7, 1.8, & 1.9!] 

LINERS [low-ionization nuclear emission line regions] 

LLAGN [low-luminosity AGN] 

Radio-loud AGN 
Radio-quiet AGN 

BLAZERS, consisting of BL Lacs and OVVs [optically violent variables] 

Flat spectrum radio quasars (FSRQs) 

Steep spectrum radio quasars (SSRQs) 

Starbursts 

QSOs and Quasars, these two being the most extreme types of AGNs 
Phew! 

So what is really going on? 

The answer is quite simple. At the center of the galaxy is a supermassive black 
hole surrounded by an accretion disc. This disc is very hot toward the black 
hole but cooler farther out. Theoretical studies suggest that the inner part of 
the disc can be very narrow, and that the black hole is hidden deep within this 
narrow central area. The outer part of the disc is believed to be a large, dense 
torus-shaped object consisting of dusty gas. Material falls into the black hole via 
the accretion disc, and prestigious amounts of energy are emitted. Sometimes 
this energy is focused into jets, as in Messier 87, which we can sometimes see. 
Moreover, this energy causes nearby clouds of fast-moving hydrogen to emit 
very strongly in the hydrogen alpha wavelength; other circumstances give rise to 
clouds farther away from the center emitting light, too. 




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Astrophysics is Easy 


Broadly speaking, the type of active galaxy (or AGN) that we see depends on 
how the accretion disc is inclined to our line of sight. If we can see the inner 
regions, it could be a Seyfert 1 active galaxy. If the inner region is obscured from 
view, it may be a Seyfert type 2. It is important to realize that even though we 
may see a galaxy face-on, it does not necessarily mean we can see the active 
nucleus face-on. The accretion disc may be inclined at a very steep angle to the 
plane of the galaxy, as in Centaurus A. Figure 4.2 attempts to demonstrate this. 

As I mentioned earlier, the source for all the energy in active galaxies is believed 
to be the supermassive black holes that lurk at their centers. Further observations 
suggest that most active galaxies are interacting or merging with nearby smaller 
galaxies, and these events provide a source of material than can feed the central 
supermassive black hole. This leads nicely to the idea that galaxies that are not 
interacting with a companion, or have not done so in the recent past, will not 
have material flowing into the black hole, and thus will not be active. And indeed, 
this is what we see. 

A certain type of AGN that has profound consequences on both galaxy 
evolution and cosmology are the Quasars, also known as quasi-stellar objects, or 
QSOs. It is important that we briefly discuss them here. The story begins in 1963, 


Blazar 


Broad 



Seyfert 1 


Figure 4.2. Diagrammatical form of the unified theory of active galactic nuclei. 



Galaxies 


179 


when Maarten Schmidt at Hale Observatories managed to identify some previ- 
ously unknown spectral lines in a supposedly stellar spectrum. He discovered 
that the unknown lines were in fact hydrogen lines, but with large redshift. For 
the object he was looking at, 3C 273, he measured a redshift of 15.8%. This is not 
particularly large as redshifts go, but soon after, quasars with larger redshifts were 
discovered. “So what?”, you may ask! Well, the significance is that these quasars 
are very far from us, as evident from their redshifts; however, they appear bright 
in photographs, in fact, like stars, and hence the quasi-stellar nomenclature. To 
be so easily seen, 14 and yet at immense distances from us, quasars must be very 
luminous, perhaps as much as 10 to 1000 times as luminous as a galaxy. Thus, 
quasars must be superluminous. 

Another important factor appears now; some quasars fluctuate in brightness 
in only a few months, and since an object cannot alter its brightness appreciably 
in less time than it takes light to cross its diameter, then these quasars must be 
very small objects, perhaps not more than a few light-months in diameter. What 
can possibly make quasars have nearly 1000 times more energy than all the stars 
in a galaxy, yet be so small? You guessed it — a supermassive black hole. 

We now know that quasars are the nuclei of galaxies that lie at tremendous 
distances from us, and thus must be objects that formed in the early universe. 
Deep images suggest that quasars contain supermassive black holes, and these 
young galaxies are distorted, and many have close companions. The activity we 
observe is possibly initiated by interactions between the host and companion 
galaxies. 

It is a nice aspect of observational astrophysics. Several active galaxies can 
be easily observed by amateur astronomers, as the brief list that follows would 
indicate. 

4.7.1 Brightest Active Galaxies 

Caldwell 29 NGC 5005 13 h 10.9 m +37° 03' April 9 

9.8m [12.6m] 6.3' 1 3.0' SAB(rs)bc easy 

This galaxy is not a binocular object, and so telescopes of about 15 cm aperture 
will reveal it only as an oval patch with a bright nucleus. The galaxy does not 
have any conspicuous spiral arms, and so even with large-aperture telescopes 
further detail will be sparse, and only a slight irregularity in overall brightness 
will be resolved. Although it is similar to the Milky Way, what makes this galaxy 
special is that it is an active galaxy of a class called LINER (Low Ionization 
Nuclear Emission-line .Region). At the center of the galaxy is some sort of 
mechanism that gives rise to both the observed spectral lines and a radio source. 
It may be due to massive stars called warmers, or an accretion disc around 
a black hole. 

Messier 77 NGC 1068 02 h 42.7 m -00°01' November 1 

8.9m [13.2 m] 7.1' |6.0' ~ (R)SA(rs)b easy 

This is a famous galaxy for several reasons. With binoculars, it is visible just 
as a hazy patch of light; under excellent seeing conditions, a faint star-like 


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Astrophysics is Easy 


nucleus maybe glimpsed. With telescopes of aperture about 10 cm or greater, and 
providing that dark skies are available, the spiral arms can be glimpsed. What 
makes this galaxy so special is that it is the archetypal active galaxy of a class 
known as Seyferts. Its uniqueness was discovered in the middle of the twentieth 
century by Carl Seyfert, who noticed that it had very prominent emission lines. 
These are due to the high velocity of gas close to the nucleus of the galaxy. The 
high speed of the gas, in the order of 350 kilometers per second, is believed to be 
due to the influence of a massive black hole. M77 is in fact classified as a Seyfert 
II galaxy, which indicates that it has only narrow emission lines. Seyferts are 
distant cousins of the famous quasars. It is one of the brightest active galaxies 
visible even to amateur astronomers. 

Caldwell 67 NGC 1097 02 h 46.3 m -30° 16' November 2 

9.5m [13. 6 m] 9.3'|6.6' w>“ SB(s)b easy 

This is a nice galaxy whose bar can be seen easily. With a 20 cm-aperture 
telescope, the core can be resolved, and a faint elongated glow is easy to see, 
which in fact is the bar. Larger apertures will resolve this feature quite well, along 
with the spiral arms that emanate from the bar’s end. It is an active galaxy, and 
classified as a Seyfert galaxy of type 1. This means that gas close to the nucleus is 
moving at extremely fast speeds, maybe in excess of 1000 kilometers per second. 
The most likely cause of this motion is the influence of a supermassive black 
hole. A Seyfert 1 galaxy has both broad and narrow emission lines, the width of 
the line being a measure of the velocity of the gas that produced the emission 
line. 

Messier 87 NGC 4486 12 h 30.8 m +12°24' March 30 

8.6m [12.7 m] 8.3'|6.6' E0.5P easy 

A very special galaxy, M87 is bright and easily seen with binoculars, and with 
telescopes a little more can be resolved. But its rather bland appearance is 
deceiving. This is a monster galaxy with a mass estimated to be that of 800 
billion Suns. This makes it one of the most massive galaxies known in the entire 
universe. But that is not all. It is an active galaxy, and lurking at its core is 
a supermassive black hole with a mass of 3 billion Suns. Another feature that 
some observers report seeing with telescopes of aperture 50 cm is the famous 
“jet” that streams out from M87. It would be a challenge indeed, and a triumph, 
if it were ever observed from the light-polluted skies of the UK. The jet is a 
stream of plasma (hot ionized gas), several thousand light years in length, which 
is believed to be due to some sort of interaction between the black hole and its 
surroundings. It is, however, very easy to photograph and image with a CCD 
camera. M87 lies at the heart of the Virgo Cluster, and most of the surrounding 
galaxies are influenced by its tremendous gravitational attraction. The cluster has 
about 300 large galaxies and perhaps as many as two thousand smaller ones. It 
is the closest large cluster, lying at a distance of around 55 million l.y. It spans 
over 100 square degrees in both Virgo and Coma Berenices. Such is its influence 
that the Milky Way is actually gravitationally attracted to it. 

Messier 82 NGC 3034 09 h 55.8 m +69°41' February 19 

8.4m [12.8 m] 11.2'|4.3' IOsp easy 


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181 


It is a very strange galaxy that becomes readily apparent when seen through a 
telescope. It can also be glimpsed with binoculars, where it would appear as an 
elongated pale glow. Large binoculars will hint at some more detail, and with 
averted vision the dark dust lane may be seen. With even a small telescope 
of 10 cm aperture, it becomes evident that something strange has happened to 
M82. The western part is obviously brighter than the eastern part. The core 
region appears jagged and angular. Throughout the length of the galaxy, the 
starlight appears to stream through the gaps in the dark dust lanes. It is a 
galaxy that will reward long and detailed study, especially with large aperture 
and high magnification. The galaxy is an active galaxy of the starburst type, 
and is undergoing an immense amount of star formation. This may have been 
caused by the close passage of its companion M81. Nearly 40 million years ago, 
the gravitational effect of M81 may have caused the interstellar material within 
M82 to collapse and form new stars. Subsequently, the material that was dragged 
from M82 is now believed to be falling back onto it, which gives rise to both its 
appearance and the new era of star formation. Both M81 and M82 can be seen 
in the same field of view and are a stunning sight. 

Caldwell 77 NGC 5128 13 h 25.5 m -43°01' April 13 

6.8m [12.9 m\ 18.2' 1 14.5' SOpec easy 

It is also known as Centaurus A. Although this galaxy is too far south to be 
visible to UK observers, it nevertheless is spectacular. Photographs show it as 
nearly a circular object bisected by a very prominent dark dust lane. It is visible 
with binoculars as a hazy star; larger binoculars just give a glimpse of the famous 
dark lane. With small telescopes of aperture 15 cm, the dark lane can be easily 
seen. Larger apertures will of course give a more detailed view, with the dark 
lane showing some structure. It is intriguing, especially because of its status as 
the nearest active galaxy. Note that this object has a compact core that probably 
houses a supermassive black hole, although it cannot be seen directly with the 
naked eye. This makes Centaurus A, for me, one of the most exciting objects in 
the sky. 

Caldwell 24 NGC 1275 03 h 19.8 m +41°30' November 10 

11.9m [13.2 m\ 3.5' |2.5' P difficult 

The last galaxy in our list is a very important galaxy, even though it is not visually 
impressive. It should be observed for several reasons. It is the main member of 
the Perseus Galaxy Cluster, also known as Abell 426. Several observers have stated 
that it can be seen through telescopes as small as 15 cm aperture, while larger 
apertures will make it easier to locate. If you look at it through a telescope of aperture 
40 cm or larger, this galaxy will appear to be surrounded by several fainter ones, 
which are all part of the cluster. In many respects, it is the most concentrated 
field of galaxies in the winter sky for northern observers. Caldwell 24 is a strong 
Radio Galaxy and is believed to be the remnant of a merger between two older 
galaxies. It is one of more than 500 members of the Perseus Cluster. The cluster 
itself is part of an even larger supercluster, the Pisces-Perseus Supercluster. 

3C 273 12 h 29.1 m 02°03' March 29 

12.8m Redshift (z) 0.158 2,000,000,000 l.y. 


Astrophysics is Easy 


182 
J7 

This quasar is the brightest in the sky and within the reach of medium- and 
large-aperture telescopes. There are reports that this quasar has been glimpsed 
with telescopes of 20 cm, and thus is well within the reach of amateurs. Averted 
vision will also help locate this distant object. To locate the quasar, the following 
directions should help: 

It is situated about 3.5° northeast of Eta Virginis. 

Locate the galaxy NGC 4536 (magnitude 10.6, surface brightness 13.2, at 
position R.A. 12 h 34.5 m Dec. 02° 11'). At about 1.25° east of this galaxy is the 
quasar. 

In its immediate vicinity is a double star, arranged east-west with 3 arcsecs 
separation. The double system has magnitudes 12.8 and 13, and the quasar is a 
bright, blue-tinted stellar object east of the double system. 

PKS 405-123 MSH 04-12 04 h 07.8 m -12° 11' November 22 

14.8m Redshift (z) 0.57 6,000,000,000 l.y. 

This is another quasar that should be within the reach of amateur astronomers, 
and it can be glimpsed with a telescope of 20 cm. It is located in the constellation 
Eridanus. The quasar lies about 3° to the northeast of Zaurak (Gamma Eri). When 
seen through an eyepiece, you may spot a tiny green dot to the left. This is the 
planetary nebula NGC 1535 (Cleopatra’s Eye). If you managed to see the quasar, 
and you will need detailed star maps to confirm your observations, then you are 
a member of a very small and elite group of observers. It is also incredible to 
note that the light that enters your eye from this quasar started its journey some 
1.5 billion years before the Solar System was formed! 


4.8 Gravitational Lensing 


Before I leave the topic of quasars, I should mention that it is possible, under the 
right conditions, to see one of the most fascinating consequences of Einstein’s 
general theory of relativity — gravitational lensing. 

The general theory of relativity is far beyond the scope of this book, but a 
simple outline of it is fairly sufficient. Gravity has the ability to “bend” light, if 
the gravitational force is strong enough. The first experimental justification of 
Einstein’s theory was in fact a measure of this light bending. On May 29, 1919, 
the British astronomer Arthur S. Eddington measured the amount of starlight 
was deflected by the Sun. He used a total eclipse of the Sun so that any faint 
stars would not be rendered invisible by the glare of the Sun. The accuracy of the 
measurements was about 20%, but it was enough to vindicate the theory. Subse- 
quent measurements using radio waves have managed to confirm the predictions 
made by Einstein to within 1%. 

The Sun is not the only object that can bend a ray of light. Any object that has 
sufficient mass can deflect light waves. Calculations show that when light rays 
from a distant object pass close to a compact but massive galaxy, the bending 
of the light can result in the appearance of multiple or twisted images. It is as if 
the galaxy were acting like a lens, and so any object emitting light from behind 




Galaxies 


183 


the galaxy has its light bent as it passes close to the galaxy. This bizarre effect is 
called gravitational leasing. 

In 1979, some astronomers noticed that a pair of quasars known as QSO 
0957 + 561 had identical spectra and redshifts, and it was suggested that these 
two quasars may in fact be one, and that the two images were produced by an 
intervening object. This was subsequently proven to be the correct explanation, 
whereby the light from distant quasars was being lensed by an intervening cluster 
of galaxies. 

You may have seen images of such objects in various books and magazines, 
and may have thought that it would be nearly impossible to see these through a 
telescope. The examples given are usually of quasars that are so distant that the 
Hubble Space Telescope or at least the world’s largest ground-based telescopes 
are needed to image them. 

This is (only) more or less true, but one or two quasars can be and have 
been seen by amateur astronomers. It is not easy. I must stress that good seeing 
conditions are essential to observe these faint objects, and that a detailed star 
atlas is required to confirm the observations. 

Twin quasar Q0957+0561A/B 10 h 01 m 55°53' February 20 

16.8m (17.1 17.4 A/B) Separation 6" 8,000,000,000 light years 

The quasar is in the constellation Ursa Major, and so it is a fine target for 
observers from the northern hemisphere. The starting point for the quasar is the 
bright edge-on galaxy NGC 3079 (8.1'x 1.4', magnitude 11.5, within the reach of 
a 20 cm telescope; several fainter galaxies lie nearby). The galaxy points to the 
quasar to the southeast, about two galaxy lengths away near a parallelogram of 
13th- and 14th-magnitude stars. The quasar lies off the southeast corner. The 
two components are 17.1 and 17.4 magnitude, separated by 6". Observers with 
very large instruments of aperture 50 cm have reported seeing the two objects 
cleanly split. Like most quasars, Q0957+0561 is slightly variable in brightness. 
With small telescopes, the two images will appear as one, but slightly elongated. 
In this case, the lensing is done by a cluster of galaxies, which lie 3.5 billion l.y. 
away, and is splitting the light of the more distant Q0957+0561 into multiple 
images. Two of these images are much brighter than the others, and this is what 
was observed. It may be wise to try as high a magnification as possible. This is a 
good observing challenge for CCD users. The quasar lies at a distance of almost 
8 billion l.y. and may well be the most distant object visible to the amateur 
astronomer. 

Leo Double Quasar QSO 1120+019 ll h 23.3 m 01°37' March 13 

15.7 20.1m (A/B) Redshift (z) 1.477 

This is an extremely difficult quasar to resolve. The brighter A component 
can be easily seen with large telescopes, but the fainter B component is very 
difficult. 


Cloverleaf Quasar H 1413+117 14 h 15.8 m 11°29' April 25 

17m(A/B/C/D) Redshift (z) 2.558 


184 


Astrophysics is Easy 


It is an exceedingly difficult object to observe, except under perfect conditions 
and with very large-aperture telescopes. The greatest separation among the four 
is about 1.36". To my knowledge, it has never been observed from the UK, but 
US observers report seeing just an asymmetric, faint hazy, and tiny blob of light, 
although it has been imaged by CCD. 


4.9 Redshift, Distance, 
and the Hubble Law 


During the early part of the 20th century, astronomers, notably Edwin Hubble, 
were beginning to take accurate measurements of the spectra of galaxies. What 
soon became apparent was that many of them exhibited a redshift, that is, the 
spectral lines were shifted to the red end of the spectrum, indicating that the 
galaxy was moving away from the Earth. Furthermore, as astronomers such 
as Milton Humason, who refined the distance determination measurements, 
observed, galaxies that were farthest away seemed to have the greatest redshifts. 
This phenomenon was observed not in just a few galaxies but in thousands, and 
as more and more measurements were taken, and as techniques improved, it 
became apparent that throughout the observed universe galaxies were moving 
away from us, and those that were far away moved faster toward the most distant 
galaxies with the greatest velocity. 

What had been discovered was the expansion of the universe and the 
Hubble Law. 

Although a rigorous treatment of this topic is beyond the scope of this book, 
it is very easy [within limits] to determine both the redshift and recessional 
velocity of a galaxy. As an aside, it is important to mention that to be completely 
accurate it is really galaxy clusters, and thus the galaxies contained therein, that 
are moving away from each other. Galaxies, say within the Local Group, actually 
have random motions. Understand the fact that we are moving toward M31 and 
that the Large Magellanic Cloud is moving toward us! 

The Hubble Law is one of the most important concepts in astrophysics, as not 
only does it relate to distance and velocity but it has more profound implications. 
It deals with the subject of cosmology and begins to set the scene for the greatest 
of all concepts — the Big Bang. 


Box 4.1 : Redshift 

The redshift of an object is the difference between the observed wavelength of a spectral 
line and its rest wavelength. 

Redshift = z = A ° faer ^~ Are5t 

•\-rpct 





Galaxies 


185 


Example: 

A galaxy has an observed Ha line at 662.9 nm. The rest wavelength of Ha is 656.3 nm. 
Calculate the redshift of the galaxy and its velocity of recession. 

662.9-656.3 

z = = 0.010 

656.3 

The redshift of the galaxy is 0.010 

For nearby galaxies, when z is much less than 1, its velocity can be calculated by 

v = c x z 

where c — speed of light, 3.0 x 10 s m/s. Thus 

v = c x z = (3.0 x 10 s m/s) x 0.01 = 3, OOOkm/s 
The galaxy’s velocity of recession is 3,000 km/s. 


Hubble's Law 


Using Hubble’s Law, we can determine the distance to a galaxy if we know how fast it 
is moving away from us. This velocity is called the recessional velocity. 

v = H 0 x d 


H 0 is called the Hubble constant, and it is generally quoted in kilometers per second 
per megaparsec (km/s/Mpc). The value [currently!] appears to be about 70km/s/Mpc. 

Example: 

Estimate the distance to the galaxy used above. 



3,000 

70 


43Mpc 


The galaxy is approximately 43 Mpc away. As 1 Mpc = 3.26 million l.y., this is 
equivalent to 140 million l.y. 


4. 1 0 Clusters of Galaxies 


Surprisingly, single galaxies are a rare breed. Most galaxies live in clusters, which 
may contain just a few; some are giant clusters, which may have thousands of 
members. In addition, small clusters occupy only a relatively small region of 







186 


Astrophysics is Easy 


space, say 1 Mpc, while the largest clusters cover an immense lOMpc. The Milky 
Way is a member of a small cluster called the Local Group, with more than three 
dozen 15 other galaxies. 

One can consider two types of clusters: Rich Clusters and Poor Clusters. The 
former may consist of more than 1000 galaxies, lots of ellipticals, and cover an 
area of over 3 Mpc in diameter. In this type of cluster, the galaxies are more 
often concentrated towards the cluster center. At the center itself, there may be 
one or two giant elliptical galaxies. A fine example is the Virgo cluster, with the 
giant elliptical M87 at its center. 

Poor clusters, as the name suggests, contain fewer than a thousand members 
that cover an area as big as that of a rich cluster, thus making the galaxies more 
spread out. 

It seems that rich clusters contain about 80-90% E-type and SO-type galaxies, 
with a few spirals, whereas poor clusters have a larger proportion of spiral 
galaxies. Furthermore, for galaxies that are in isolation (i.e., those not in 
clusters), it appears that 80-90% of these are spirals. There is a large amount 
of evidence suggest that large elliptical galaxies have been involved in to many 
galaxy collisions, whereas spirals have not. In fact, it may well be that ellip- 
ticals are formed by the merger of spirals. The dwarf elliptical, on the other 
hand, seems to have followed a different evolutionary path. These are small 
galaxies that have lost their gas and dust due to interactions with several larger 
galaxies. 

Although the evolution of galaxies is still not fully understood, it is obvious 
that interactions between them is very important. Collisions, mergers, and close 
encounters can all cause bursts of rapid star formation and very dramatic and 
spectacular tidal disruption. In fact, our own Milky Way, I think, is cannibalizing 
the Magellanic Clouds. 

Amazing! 


4. 1 0. 1 Groups and Clusters of Galaxies 

Hickson Group 68 NGC 5353 13 h 53.4 m +40°47' April 20 

11.1m — > 11.2'-<— 5 moderate 

This is a very nice group of stars for amateur instruments. The brightest member 
can be seen with a telescope as small as 6 cm, and with a 15 cm aperture, it will 
show a slight brightening at its center. The other galaxies will appear as faint 
patches of light; to see the faintest member would most certainly require an 
aperture of about 25 cm. 

Copelands Septet NGC 3753 ll h 37.9 m +21°59' March 16 

13.4m — ► 7.0 ' ■*— 7 difficult 

Also known as the Hickson Galaxy Group 57, this is a very small group of 
galaxies. It is situated in the constellation Leo, all within about 7 arcsecs of each 
other. Even telescopes of 25 cm aperture will not spot the fainter members of 


Galaxies 



the group, but just the four brighter galaxies. Larger apertures should, of course, 
help spot them. Nevertheless, seeing conditions will determine what you observe. 
The group is a mix of barred spirals, ordinary spirals, and lenticular galaxies. 


This is a large cluster of galaxies. Many of its members are within the reach of 
amateur telescopes of aperture 25 cm or larger. It is fairly well spread out, and 
so the field of view will be dotted with many indistinct faint patches of light. It 
contains many elliptical, spiral, barred spiral, and lenticular galaxies. 


It is a very famous group of galaxies located in Pegasus, but one that was proved 
strangely difficult for amateurs in the past. Under perfect seeing conditions, the 
group is visible with a 20 cm telescope. However, I stress the word perfect ! The 
largest member of the quintet is only 2.2 x 1.2 arcsecs in size, and so it is very 
small but the brightest. To observe the group as a distinct unit, and not a faint 
smudge of light, you will need a telescope of aperture 25 cm. This will show 
at least four of the group, but the fifth one would require an aperture of at 
least 30 cm. With high magnification and a larger aperture, the structure can be 
seen within the brighter members. It is believed that the four in the group are 
interacting with each other. There has been debate as to whether the fifth is in 
fact a line-of-sight galaxy. Finally, it is a definite observing challenge from an 
urban locality. 


This is a real challenge! Except with the largest telescopes, it is doubtful that you 
can see anything at all, and even with apertures around 40 cm, the galaxies will 
barely be resolved. Nevertheless, it would be interesting to find out what would 
be the smallest aperture required to spot these faint galaxies. 

Fornax Cluster NGC 1316 03 h 20.9 m -37°17' November 10 
11.4m — ► + 12'-*— 10+ difficult 

This is another large cluster of galaxies. Amateur telescopes should be able 
to pick out the brightest members with no difficulty. What makes this cluster 
spectacular, however, is that, with a modest aperture (say 25 cm) and clear dark 
skies, there are so many galaxies visible that identifying it will be very difficult. 
The brightest member is visible even with an 8 cm telescope! A galaxy of note 
in the cluster is NGC 1365, which is a nice barred spiral 8 arcsecs in length and 
visible as a faint blur in an 8 cm aperture telescope. The cluster also contains a 
galaxy known as the Fornax System. It is a dwarf spheroidal galaxy, a very small 
and faint class of galaxy. 


Coma Cluster NGC 4889 12 h 57.7 m 

11.4m — ► 120+'-<— 10+ 


+28° 15' April 6 
difficult 


Stephen's Quintet NGC 7320 22 h 36.1 m 

12.6m — ►4'-<— 5 


+33° 57' April 6 
difficult 


Seyferf s Sextet NGC 6027 15 h 59.2 m 

13.3m 1.5' <— 6 


+20° 46' May 22 

very difficult 


188 


Astrophysics is Easy 


4.11 Endnote 


We have now come to the end of our spectacular journey, and I hope you have 
enjoyed the trip as well as being amazed and sometimes astounded by what you 
have read and hopefully observed. But this is not the end — it is just the beginning 
because you have only seen a handful of the plethora of celestial delights that 
await you. 

The next time you observe the night sky, just think: you will have an inkling 
of what those objects are, how they formed, how they could die, what they are 
made of, and whether they are stars, clusters, nebulae, or galaxies. 

Incredible! 

Happy observing! 


Notes 


1. The Milky Way galaxy is often referred to as the “Galaxy,” with a capital 
letter, whereas any other is simply a “galaxy.” 

2. Although a few stars may, after an immense amount of time, break free of 
a galaxy’s grip and become intergalactic wanderers. 

3. Or, at times, a rugby ball (for those of us in the UK). 

4. See Chapter 2 for a description of HII regions. 

5. Recent research has shown that the Large Magellanic Cloud, often classified 
as irregular, is in fact a spiral galaxy, even though it bears little resemblance 
to the classic spiral shape. 

6. Of course, I don’t really have to mention that if you have a medium-to- 
large-aperture telescope, then the number of galaxies visible to you is vast, 
and the detail you will be able to see will astound you! 

7. The Local Group is a cluster of several galaxies, including the Milky Way. 
It consists of M31, M33, M110, and M32, the Large and Small Magellanic 
Clouds, and about 25 other dwarf galaxies, including Leo I and II, And I 
and II, the Draco, Carina, Sextans, and Phoenix dwarfs. 

8. Cepheid variables are used as standard candles, which measure distances 
to other extra-galactic objects. 

9. The galaxy M83 lies at the same distance and has reportedly been seen with 
the naked eye. 

10. Sometimes the cluster is just referred to as the Virgo Cluster. For a 
description of the cluster, see the entry on M87. 

11. A large HII region, NGC 604, is visible. See the entry under Emission 
Nebulae in Chapter 2. 

12. Active galactic nuclei are usually classified by three parameters: optical 
variability, radio emission, and spectral line width. Seyfert galaxies are 
mostly radio-quiet and not so strongly variable as other types of AGN; 
Seyfert Is have broad and narrow spectral lines, while Seyfert 2s only have 
narrow lines. In addition, some Seyfert Is exhibit broadened lines that are 
relatively narrow and are designated “Narrow Line Seyfert Is” (NLSls). 




Galaxies 


189 


Quasars are broad-lined and some of them are variable; they can be further 
divided between radio-quiet and radio-loud quasars. Radio galaxies have 
strong radio emission and are not variable. Blazers are highly variable and 
some of them have strong radio emission; they can be divided between 
narrow-lined BL Lac objects (some of them show no lines of spectra at all; 
that’s why redshift of many BL Lacs remains unknown) and broad-lined 
OVVs (Optically Violently Variable quasars). 

13. And this is not a complete list, either! 

14. By easily, I mean photographically. 

15. Additional members of the Local Group are being found at regular intervals. 
These are small and indistinct, thus their difficulty in being observed! 


APPENDIX ONE 


Degeneracy 


Degeneracy is a very complex topic but a very important one, especially when 
discussing the end stages of a star’s life. It is, however, a topic that sends quivers 
of apprehension down the back of most people. It has to do with quantum 
mechanics, and that in itself is usually enough for most people to move on, and 
not learn about it. That said, it is actually quite easy to understand, providing that 
the information given is basic and not peppered throughout with mathematics. 
This is the approach I shall take. 

In most stars, the gas of which they are made up will behave like an ideal gas, 
that is, one that has a simple relationship among its temperature, pressure, and 
density. To be specific, the pressure exerted by a gas is directly proportional to 
its temperature and density. We are all familiar with this. If a gas is compressed, 
it heats up; likewise, if it expands, it cools down. This also happens inside a star. 
As the temperature rises, the core regions expand and cool, and so it can be 
thought of as a safety valve. 

However, in order for certain reactions to take place inside a star, the core 
is compressed to very high limits, which allows very high temperatures to be 
achieved. These high temperatures are necessary in order for, say, helium nuclear 
reactions to take place. At such high temperatures, the atoms are ionized so that 
it becomes a soup of atomic nuclei and electrons. 

Inside stars, especially those whose density is approaching very high values, 
say, a white dwarf star or the core of a red giant, the electrons that make up the 
central regions of the star will resist any further compression and themselves set 
up a powerful pressure . 1 This is termed degeneracy, so that in a low-mass red 


191 


192 


Astrophysics is Easy 


giant star, for instance, the electrons are degenerate, and the core is supported 
by an electron-degenerate pressure. But a consequence of this degeneracy is that 
the behavior of the gas is not at all like an ideal gas. In a degenerate gas, the 
electron degenerate pressure is not affected by an increase in temperature, and 
in a red giant star, as the temperature increases, the pressure does not, and the 
core does not expand as it would if it were in an ideal gas. The temperature, 
therefore, continues to increase, and further nuclear reactions can take place. 

There comes a point, however, when the temperatures are so high that the 
electrons in the central core regions are no longer degenerate, and the gas behaves 
once again like an ideal gas. 

Neutrons can also become degenerate, but this occurs only in neutron stars. 

For a fuller and more rigorous description of degeneracy, I recommend the 
books mentioned in the latter appendices. Be warned, however, that mathematics 
is used liberally. 


Note 


1. This is a consequence of the Pauli exclusion principle, which states that two 
electrons cannot occupy the same quantum state. Enough said, I think! 


APPENDIX TWO 


Books, Magazines, 
and Astronomical 
Organizations 


Books, Magazines, 
and Organizations 


There are many fine astronomy and astrophysics books in print, and to choose 
among them is a difficult task. Nevertheless, I have selected a few, which I believe 
are among the best on offer. I do not expect you to buy, or even read, them all, 
but it would be in your better interest to check at your local library to see if they 
have some of them. 


Star Atlases and Observing 
Guides 


Norton’s Star Atlas and Reference Handbook, 20th edn, Ian Ridpath (ed.), Prentice Hall, 
2003 USA. 

Sky Atlas 2000.0, W. Tirion & R. Sinnott, Sky Publishing & Cambridge University Press, 
1999, Massachusetts, USA. 

Millennium Star Atlas, R. Sinnott & M. Perryman, Sky Publishing, 2006, Massachusetts, 
USA 

Uranometria 2000.0, Vols. 1 & 2, Wil Tirion (ed.), Willmann-Bell, 2001, Virginia, USA. 


193 



1 94 Astrophysics is Easy 

Observing Handbook and Catalogue of Deep-Sky Objects, C. Luginbuhl & B. Skiff, 
Cambridge University Press, 1990, Cambridge, UK. 

The Night Sky Observer’s Guide, Vols. I & II, G. Kepple & G. Sanner, Willman-Bell, 1999, 
Richmond, USA. 

Deep-Sky Companions: The Messier Objects, S. O’Meara, Cambridge University Press, 1999, 
Cambridge, UK. 

Deep-Sky Companions: The Caldwell Objects, S. O’Meara, Cambridge University Press, 
2004, Cambridge, UK. 

Observing the Caldwell Objects, D. Ratledge, Springer-Verlag, 2000, London, UK. 

Burnham’s Celestial Handbook, R. Burnham, Dover Books, 1978, New York, USA. 

Star Clusters and How To Observe Them, M. Allison, Springer, 2006, New York, USA. 

Double & Multiple Stars and How To Observe Them, J. Mullaney, Springer, 2006, 
New York, USA. 

Nebulae and How To Observe Them, S. Coe, Springer, 2006, New York, USA. 

Galaxies and How To Observe Them, W. Steinicke & R. Jakiel, Springer, 2006, New York, 
USA. 


Astronomy and Astrophysics 
Books 


Astrophysical Techniques, 4th edn, C. Kitchin, Institute of Physics, 2003, Bristol, UK. 

Discovering the Cosmos, R. Bless, University Science Books, 1996, Sausilito, USA. 

The Cosmic Perspective, J. Bennett, M. Donahue, N. Schneider, M. Voit, Addison Wesley, 
1999, Massachusetts, USA. 

Voyages Through The Universe, A. Fraknoi, D. Morrison, S. Wolff, Saunders College 
Publishing, 2000, Philadelphia, USA. 

Introductory Astronomy & Astrophysics, M. Zeilik, S. Gregory, E. Smith, Saunders College 
Publishing, 1999, Philadelphia, USA. 

Pathways To Astronomy, Schneider & Arny, McGraw-Hill, 2006, USA. 

An Introduction To The Sun & Stars, Green & Jones, Open University/Cambridge 
University Press, 2005, Cambridge, UK. 

An Introduction To Galaxies & Cosmology, Jones & Lambourne, Open University/ 
Cambridge University Press, 2005, Cambridge, UK. 

Introduction to Modern Astrophysics, 2nd edn, B. W. Carroll & D. A. Ostlie, 
Addison Wesley, 2006, USA. 

Stars, J. B. Kaler, Scientific American Library, 1998, New York, USA. 

Extreme Stars, J. B. Kaler, Cambridge University Press, 2001, Cambridge, UK. 

The Physics of Stars, 2nd edn, A. Phillips, Wiley, 1999, Chichester, UK. 

Stars, Nebulae and the Interstellar Medium, C. Kitchin, Adam Hilger, 1987, Bristol, UK. 

100 Billion Stars, R. Kippenhahn, Princeton University Press, 1993, Princeton, USA. 

Stellar Evolution, A. Harpaz, A. K. Peters, Ltd, 1994, Massachusetts, USA. 

The Fullness of Space, G. Wynn-Williams, Cambridge University Press, 1992, Cambridge, 
UK. 

Astrophysics Of Gaseous Nebulae And Active Galactic Nuclei, 2nd edn, D. E. Osterbrock & 
G. J. Ferland, University Science Books, 2005, Sausilito, USA. 

The Dusty Universe, A. Evans, John Wiley, 1994, Chichester, UK. 

Galaxies and Galactic Structure, D. Elmegreen, Prentice Hall, 1998, USA. 


Books, Magazines, and Astronomical Organizations 


195 


Exploring Black Holes, E. Taylor & J. A. Wheeler, Princeton University Press, 2001, 
Princeton, USA. 


Maqazines 


Astronomy Now, UK 
Sky & Telescope, USA 
New Scientist, UK 
Scientific American, USA 
Science, USA 
Nature, UK 

The first three magazines are aimed at a general audience and so are applicable 
to everyone, while the last three are aimed at the well-informed lay person. In 
addition, there are many research-level journals that can be found in university 
libraries and observatories. 


Organizations 


The Federation of Astronomical Societies, 10 Gian y Llyn, North Cornelly, 
Bridgend County Borough, CF33 4EF, Wales 
[http://www.fedastro.org.uk/ ] 

Society for Popular Astronomy, The SPA Secretary, 36 Fairway, Keyworth, 

Nottingham, NG12 5DU, UK 

[http://www.popastro.com/] 

The American Association of Amateur Astronomers, P.0. Box 7981, Dallas, 

TX 75209-0981 

[http://www.astromax.com/] 

The Astronomical League 
[http://www.astroleague.org/] 

The British Astronomical Association, Burlington House, Piccadilly, London, 
W1V 9AG, UK 

[http://www.britastro.org/baa/] 

The Royal Astronomical Society, Burlington House, Piccadilly, London W1V ONL 
[http://www.ras.org.uk/membership.htm] 

International Dark Sky Association 
[ http :// www. darksky.org/] 





Topic Index 


Absolute magnitude, 4, 10, 11, 35, 37, 117 
Absorption lines, 23-27, 30, 39, 44, 102, 146 
Active galactic nuclei, 177, 188 
Active galaxies, 177-180 
AGN, 177-178, 188 
Alpha particles, 105, 142 
Apparent brightness, 3, 6, 7, 8, 10, 42, 43, 109 
Apparent magnitude, 4, 9-11, 14, 43, 117, 120 
Arc-second, 1-3, 5, 6, 12, 14, 49 
Astrometric binary, 92, 154 
Astronomical unit, 1, 95, 97, 140 
Asymptotic giant branch, 109, 122-123 

Balmer lines, 26 

Barnard objects, 58, 62 

Barred spiral, 158, 170, 172, 187 

B associations, 82 

Binary stars, 92-94, 96, 155 

Bipolar outflow, 71 

Black holes, 135, 147, 149-151, 156, 178, 179 

Blazers, 177, 189 

BL Lacs, 177, 189 

Bok globules, 58 

Brightness ratio, 10 

Bulge, 108, 155, 158-160, 162, 171 

Carbon burning, 138-139, 146-147 
Chandrasekhar limit, 134-135, 137-138, 142, 
146, 149 

Circumstellar accretion disc, 71 


CNO cycle, 154-155 

Color index, 17, 127 

Color-magnitude diagram, 109 

Convection, 59, 66, 67, 86, 87, 118, 124 

Convection zone, 87 

Core, 87, 91 

Core bounce, 142 

Core collapse, 142, 149, 150, 155 

Core helium burning, 105, 107, 109, 118, 123 

Core hydrogen burning, 98, 122, 124 

Core rebound, 142 

Dark nebuke, 53-55, 58, 62, 69, 72, 79, 113 
Degeneracy, 105, 134, 138, 149-150, 155, 
191-192 

Deuterium, 89, 90 

Disc, 71-72 

Distance modulus, 11 

Doubly-ionized oxygen, 130 

Dredge-ups, 124 

Dust grains, 53, 56, 58, 129, 144 

Dwarf elliptical, 159, 174, 186 

Eclipsing binary, 22, 95, 140 

Elliptical galaxies, 154, 158-160, 172, 174, 

176, 186 

Emission lines, 23-25, 29, 70, 141, 146, 163, 180 

Emission line spectrum, 23 

Emission nebuke, 47-49, 51, 56, 83, 153, 188 


197 


198 


Topic Index 


Energy flux, 19 

Energy level, 23-26, 39, 44, 47-48, 61 
Event horizon, 150 

Evolutionary track, 63-67, 69, 102, 107-108, 
118-119, 122, 135, 139, 153 
Extended object, 162 

First dredge-up, 124 

Flat spectrum radio quasars, 177 

Fluorescence, 47, 69 

Flux, 19-20, 43 

FSRQ’s, 177 

Galactic clusters, 72 
Galactic plane, 78, 110 
Galaxies, 1, 4, 6, 58, 85-86, 146, 153-154, 
157-181, 183-189 
Gamma rays, 22, 90-91, 142 
Giant molecular clouds, 62, 73, 83 
Globular clusters, 72, 76, 108-111, 113, 119, 

123, 155 

Gravitational equilibrium, 63, 65, 88, 153 
Gravitational lensing, 182-183 
Ground state, 26 

Halo population, 5 
Helium, 22, 101 

Helium burning, 73, 104-109, 118, 122-124, 126, 
128, 138, 154-155 
Helium capture, 138 

Helium flash, 104-107, 109, 119, 122-123, 128, 
147, 155 

Helium-shell flash, 128 
Herbig-Haro objects, 71 
Hertzsprung-Russell diagram, 35 
HII regions, 47, 48, 62, 85, 159, 161-162, 172, 
176, 188 

Horizontal branch, 109, 119, 123 
Horizontal branch stars, 109 
Hubble classification, 159, 160, 163-166 
Hubble law, 4, 184 
Hubble tuning fork, 161 
Hydrogen, 22 

Hydrogen burning, 41, 67, 73, 86, 98, 101-102, 
104, 106, 122-124, 128, 138, 154 
Hydrogen burning shell, 104, 106-107, 123, 
128-129 

Hydrostatic equilibrium, 63, 66, 88, 107, 115 

Instability strip, 116, 118, 119 
Integrated magnitude, 162, 170 
Interstellar extinction, 56 
Interstellar medium, 45-62, 85, 126, 143, 145, 
149, 158, 159 
Inverse square law, 7 
Ionization, 26, 48 
Irregular galaxies, 158-159 
Isotope, 89, 105, 139 


Jeans criteria, 59-60 
Jeans length, 60-61 
Jeans mass, 60-61 

Kepler’s law, 96-97, 151, 154 

Later-type, 25 

Lenticular, 158-160, 175, 187 
LINER, 177, 179 
Lithium, 70 
LLAGN, 177 

Local group, 159, 163, 169, 175, 177, 184, 186, 
188-189 

Low-mass stars, 67, 72, 98, 106-107, 122, 128, 
133, 138, 147 

Luminosity, 3, 6-8, 12, 13, 19-22, 26, 29-33, 35, 
37, 39, 41-43, 56, 63-67, 70, 73, 79, 86, 88, 
91, 98, 102, 106-109, 112, 114, 117, 118, 
122-123, 126, 128-129, 135-136, 139, 142, 
144, 148, 153, 155, 177 
Luminosity distance formula, 7 
Lyman alpha, 48 

Main sequence, 37 

Main-sequence lifetime, 98-100, 106-107, 

109, 123 

Main-sequence star, 32, 37, 39, 42, 63-64, 66, 
69-70, 73-74, 83, 97-99, 101-102, 107-109, 
122, 130, 134-137, 155 
Mass ejection, 133 

Mass-luminosity relationship, 67, 153 
Molecular clouds, 53, 57-58, 62, 69, 71, 73, 83, 
85, 155 

Multispin, 166 
Neon burning, 139 

Neutrino, 89-90, 142-143, 147, 154-155 
Neutronization, 142 

Neutrons, 88, 90, 139, 141-142, 147, 192 
Neutron star, 15, 135, 142, 144-145, 147-149, 
151, 156, 192 

Nuclear burning, 138-139, 141 
Nuclear fusion, 22, 37, 39, 63, 65-67, 69, 73, 
85-86, 88-90, 101, 104, 107, 115, 123, 133, 
135, 137, 139 
Nucleon, 141 

OB associations, 82-83, 86 
Opacity, 54, 66, 88, 115-117, 124 
Optical doubles, 92 
OVV, 177, 189 

Oxygen, 48, 105, 123, 126, 128-130, 133, 135, 
137-139, 142, 155 

Parallax, 1-3, 5, 10, 95 
Parsec, 2, 42 

Period-luminosity relationship, 3, 42, 118 


Topic Index 


199 


Photometry, 7 
Photosphere, 87 

Planetary nebulae, 47, 122, 128-132, 136, 
141, 155 

Plasma, 87, 154, 180 

Plerion, 145, 149 

Poor cluster, 186 

Population I, 73, 118, 155, 159 

Population II, 96, 118, 119, 121, 155, 159 

Position angle, 93, 154 

Positron, 89-90 

Primary star, 6, 21, 32, 92-94, 137-138 
Proper motion, 5, 6, 13, 30, 43, 96, 138 
Proto-galaxy clouds, 109 
Proton-proton chain, 63, 88-90, 101, 153 
Protons, 48, 88-91, 139, 141-142 
Proto-stars, 51, 58-59, 63-67, 69-72, 84-85, 
107, 153 

Protostellar disk, 71 
Pulsars, 147-149 

QSO, 177-178, 183 

Quasar, 62, 177-180, 182-183, 189 

Radiation zone, 87 
Radioactive transfer, 124 
Radio loud AGN, 177 
Radio quiet AGN, 177 
Random walk, 91 
Recessional velocity, 184-185 
Red-giant branch, 122 
Redshift, 179, 181-185, 189 
Reflection nebulae, 56-57 
Rich cluster, 186 
Roche lobe, 146, 155 

Schwarzschild radius, 150-152 
Secondary star, 94, 120 
Second dredge-up, 124 
Second red-giant phase, 123 
Seyfert, 177, 180, 188 
Seyfert type I, 178, 180, 188 
Seyfert type II, 180, 188 
Shapley-Sawyer concentration 
class, 111 

Shell helium burning, 123 
Shell hydrogen burning, 101, 109, 123 
Shock waves, 85 
Silicon burning, 139 
Small molecular clouds, 73 
Solar luminosity, 106, 117 
Spectral type, 19, 25-27, 35, 37, 39, 70, 71, 
99,104, 113, 121 
Spectra of stars, 3, 23 
Spectroscopic binary, 12, 29, 31, 32, 94, 

96, 121 

Spheroidal component, 158-159 


Spiral galaxies, 154, 158-160, 162, 163, 170, 
176, 186 
SSRQs, 177 
Starburst, 177, 181 
Star formation triggers, 84 
Stars 

Ae stars, 57, 71 
AGB star, 126 
be stars, 71 
biggest, 21 
blue stragglers, 78 
brightest, 12 
brown dwarfs, 69 
carbon, 124, 126, 127-128 
Cepheid, 29, 117 
Type I, 118 
Type II, 118, 121 
clusters, 107 
galactic, 72 
globular, 108, 111 
open, 72-73 
constituents, 22 
dwarf, 38, 96, 133, 134, 137 
eclipsing binary, 22 
giants, 101, 103 
high-mass, 138-139, 147 
infrared, 125 
lifetime, 97-99 

long period variable, 19, 34, 114, 119-121 
low-mass, 72-73, 98, 106-107, 109, 122-123, 
128, 133, 138, 147 
mass, 19, 138, 147 
nearest, 5, 86 
pulsating variable, 3, 114 
RR Lyrae variable, 3, 114, 119, 121 
red dwarf, 5-6, 14, 138 
red giant, 101, 103 
subdwarf, 130 
subgiant, 26 
supergiant, 139, 140 
T Tauri, 70-71 

Type 

early, 25 
late, 25 

intermediate, 25 
white dwarf, 133-138 
Wolf-Rayet, 25, 50, 79, 131, 141 
zero-age-main-sequence, 98, 107-108 
Steep spectrum radio quasars, 177 
Stefan-Boltzmann law, 19-20, 37 
Stellar associations, 82-83, 154, 170 
Stellar classification, 14, 17, 25, 39 
Stellar parallax, 1, 5, 10 
Stellar stream, 83 

Stellar wind, 50-51, 54, 85, 125-126, 129, 

137, 146 
Sun, 86-92 


200 


Topic Index 


Supernova, 47, 54, 61, 76, 82-83, 85, 122, 126, 
141-143, 149, 155-156 
Type I, 146-147, 155 

Trumpler classification, 76 
Turnoff point, 109 

Type II, 146-147 

Supernova remnant, 47, 61, 144-145, 149, 156 
Surface temperature, 15, 17, 19-20, 26, 32, 35, 
37, 39, 42-44, 63, 65-67, 70, 98, 104, 

Vorontsoz- V ellyaminov classification 
system, 130 

106-107, 109, 116, 119, 128, 135, 153 

Wien law, 15, 17 

T associations, 83 
Thermal pulse, 129 
Transitions, 24-25, 44, 63 

X-ray binary pulsar, 148 
X-ray bursters, 148 

Triple a process, 105-106, 123, 126 

ZAMS, 98, 107 


Object Index 


119 Tauri, 140 
15 Monocerotis, 27 
2Mon, 31 
30 Ophiuchi, 113 
3C 273, 179, 181 
61 Cygni, 5, 95 

Achernar, 14, 29 
Acrux, 18 
AE Aurigae, 57 
Alcyone 29 
Aldebaran, 14, 39 
Alderamin, 31 
Algeiba, 32 
Algenib, 29, 31 
Alhena, 30 
Almach, 33 

Alpha Persei Stream, 84 
Altair, 13 
Aludra, 29 

Andromeda Galaxy, 169 
Antares, 13, 33 
Antennae Galaxies, 175 
Arcturus, 12 

b Velorum, 31 
Barnard 33, 55 
Barnard 59, 54 
Barnard 78, 55 
Barnard 86, 55 


Barnard 87, 55 
Barnard 228, 54 
Barnard 352, 55 
Barnard’s Galaxy, 176 
Barnard’s Loop, 145 
Barnard’s Star, 5 
Becklin Neugebauer 
Object, 70 
Becrux, 12 
Bellatrix, 18 
Betelgeuse, 14 
Black Eye Galaxy, 167 
Blinking Planetary, 132 
Blue Flash Nebula, 132 
Blue Snowball, 133 
Bubble Nebula, 51 
Burnham, 584, 78 

Caldwell 3, 171, 174 
Caldwell 6, 131, 132 
Caldwell 7, 163, 172 
Caldwell 11, 51 
Caldwell 13, 80 
Caldwell 14, 80 
Caldwell 15, 132 
Caldwell 17, 174 
Caldwell 18, 174 
Caldwell 19, 51 
Caldwell 20, 50 


201 


202 


Object Index 


Caldwell 21, 176 
Caldwell 22, 133 
Caldwell 24, 181 
Caldwell 26, 166 
Caldwell 27, 50 
Caldwell 29, 179 
Caldwell 30, 168 
Caldwell 31, 57 
Caldwell 32, 171 
Caldwell 33, 145 
Caldwell 34, 144 
Caldwell 35, 174 
Caldwell 36, 165 
Caldwell 38, 164 
Caldwell 39, 131 
Caldwell 40, 166 
Caldwell 41, 81 
Caldwell 43, 168 
Caldwell 44, 172 
Caldwell 46, 52 
Caldwell 48, 163 
Caldwell 49, 52 
Caldwell 51, 176 
Caldwell 52, 174 
Caldwell 53, 175 
Caldwell 54, 78 
Caldwell 55, 132 
Caldwell 57, 176 
Caldwell 59, 131 
Caldwell 60/61, 175 
Caldwell 62, 170 
Caldwell 63, 130, 132 
Caldwell 64, 77 
Caldwell 65, 169 
Caldwell 67, 180 
Caldwell 70, 169 
Caldwell 72, 172 
Caldwell 76, 78 
Caldwell 77, 181 
Canopus, 31 
Capella, 14 
Castor, 30 

Cat’s Eye Nebula, 131 
CE Tauri, 140 
Centaurus A, 181 
Cloverleaf Quasar, 183 
Cocoon Nebula, 51 
Collinder 69, 81 
Collinder 81, 82 
Collinder 316, 79 
Coma Cluster, 187 
Copelands Septet, 186 
Crab Nebula, 144-145, 149 
Crescent Nebula, 50 
CVn 94, 127 
Cygnus Loop, 144 
Cygnus X-l, 152 


Delta Leonis, 30 
Deneb, 30, 51 
Denebola, 30 
Duck Nebula, 49 
Dumbbell Nebula, 132, 133 

Eagle Nebula, 50 
Electra, 29, 81 
Enif, 33 

Epsiln Eridani, 6 
Eskimo Nebula, 131, 133 
Eta Persei, 34 
Eta Sagitai, 30 

Filamentary Nebula, 145 
Flaming Star Nebula, 57 
Fomalhaut, 13 
Fornax Cluster, 187 
FU Orionis, 71, 153 

Gacrux, 34 
Gamma Herculis, 31 
Gammak Cassiopeiae, 29 
Garnet Star, 19, 80 
Ghost of Jupiter, 131 
Gienah, 32 
Gliese 229B, 69 
Great Rift, 55, 165 
Great Sagittarius Star 
Cloud, 55 
Gum 4, 49 
GX Andromadae, 4 

Hadar, 12 
HD 7902, 80 
HD 93129A, 27 
Helix Nebula, 130133 
Hercules Cluster, 112 
Herculis, 21 
Herschel 16, 132 
Herschel 53, 133 
Hickson Group 57, 187 
Hickson Group 68, 186 
Hind’s Crimson Star, 19 
Hind’s Variable Nebula, 53, 56 
Horsehead Nebula, 56 
Hubble’s Variable Nebula, 52 
Hyades, 81, 84 
Hyades Stream, 84 

IC 405, 57 
IC 410, 57 
IC 417, 57 
IC 1396, 51, 80 
IC 2118, 145 
IC 5067, 50 
Ink Spot, 55 


Object Index 


203 


Kleinman-Low Sources, 70 
KQ Puppis, 22 

La Superba, 127 
Lacille, 6 

Lagoon Nebula, 50, 74 
Leo Double Quasar, 183 
Little Dumbbell, 133 
Local Group of Galaxies, 163 
Lynds 906, 55 

Lynds Dark Nebula 1773, 54 

Maia, 29, 81 
Merope, 18, 57 
Messier 1, 145 
Messier 3, 112 
Messier 4, 112 
Messier 5, 112 
Messier 7, 79, 133 
Messier 8, 49 
Messier 9, 113 
Messier 11, 72, 80 
Messier 13, 112 
Messier 15, 113 
Messier 16, 50, 79 
Messier 17, 50 
Messier 19, 113 
Messier 20, 49 
Messier 22, 113 
Messier 24, 79 
Messier 25, 79 
Messier 27, 132 
Messier 31, 168 
Messier 33, 169 
Messier 37, 82 
Messier 41, 77 
Messier 42, 69, 75 
Messier 44, 78 
Messier 45, 81 
Messier 48, 78 
Messier 49, 173 
Messier 51, 167 
Messier 54, 113 
Messier 57, 131 
Messier 58, 166 
Messier 59, 173 
Messier 61, 165 
Messier 63, 167 
Messier 64, 162, 167 
Messier 65, 164 
Messier 66, 164 
Messier 67, 78 
Messier 68, 111 
Messier 76, 133 
Messier 77, 179 
Messier 81, 163 
Messier 83, 168 
Messier 84, 172 


Messier 85, 176 
Messier 86, 173 
Messier 87, 177, 180 
Messier 88, 164 
Messier 89, 173 
Messier 91, 171 
Messier 92, 113 
Messier 94, 167 
Messier 95, 170 
Messier 96, 163 
Messier 97, 131 
Messier 98, 165 
Messier 99, 165 
Messier 104, 166 
Messier 106, 164 
Messier 108, 171 
Messier 109, 171 
Messier 110, 175 
Mira, 34, 121 
Mirach, 33 
Mirzim, 29 
Mu Geminorum, 140 

NGC 604, 51 

NGC 1365, 172, 187 

NGC 1435, 57, 81 

NGC 1554, 52 

NGC 2024, 52 

Norma Spiral Arm, 79 

North American Nebula, 50, 55 

Northern Coalsack, 55 

Nu Draconis, 30 

Omega Nebula, 50 
Orion Association, 83 
Orion Nebula, 50, 69 
Owl Nebula, 131 
Oyster Nebula, 133 

Parrot Nebula, 55 

Pease-1, 114 

Pelican Nebula, 50, 51 

Perseus Double Cluster, 73 

Pinwheel Galaxy, 166, 170 

Pipe Nebula (bowl), 55 

Pipe Nebula (stem), 54 

Pisces-Perseus Supercluster, 181 

PKS 405-123, 182 

Plaskett’s star, 29 

Pleiades, 76, 81 

Polaris, 31, 96 

Pollux, 12, 30 

Praesepe, 78, 84 

Procyon, 5 

Procyon B, 137 

Proxima Centauri, 5, 13 

Pup, 137 


204 


Object Index 


Q 0957+0561A/B, 183 
QSO 1120+019, 183 


TT Monocerotis, 124 
Twin Quasar, 183 


R Aqr, 103 
R Cas, 104 

R Corona Borealis, 55, 128 
R Leonis, 122, 123 
R Monocerotis, 52 
R Scl, 127 
Ras Algethi, 32, 34 
Ras Alhague, 30 
Regulus, 18 
Rigel, 14 

Rigil Kentaurus, 13 
Ring Nebula, 130, 132 
Ring-Tail Galaxies, 175 
Rosette Molecular Complex, 52 
Rosette Nebula, 52, 70 
RR Lyrae, 119, 121, 155 
RS Cyg, 103 
RT Aurigae, 120 
RV Arietis, 121 
RW Arietis, 121 

S Cephei, 127 
S Pegasi, 124 
Sadal Melik, 32 
Sadal Suud, 32 

Sagittarius Carina Spiral Arm, 50, 79 
Sagittarius Dwarf Galaxy, 113 
Saturn Nebula, 132 
Scheat, 33 

Scorpius Centaurus Association, 82, 83 

Scorpius OBI, 79 

Seyfert’s Sextet, 187 

Sharpless 2-264, 82 

Sharpless 2-276, 145 

Sirius A, 5 

Sirius B, 135, 137 

Small Sagittarius Star Cloud, 79 

Sombrero Galaxy, 166 

Spica, 12 

Spindle Galaxy, 175 
Star Queen Nebula, 50 
Stephen’s Quintet, 168 
SU Cassiopeiae, 121 
Sun, 18 

Sunflower Galaxy, 167 
Swan Nebula, 50 

T Monocerotis, 121 
T Tauri, 71, 73, 83 
T Vulpeculae, 121 

Taurus Dark Cloud Complex, 71, 81 
Tempel’s Nebula, 57 
Trapezium, 69, 85 
Trifid Nebula, 49, 56 
Trumpler, 24, 79 


U Aquilae, 121 
U Cam, 127 

Ursa Major Streamsg8, 384 
Ursae Majoris, 84, 94, 120 
UV Ceti, 6 

V Aql, 127 

V Pav, 127 

V467 Sagittari, 121 
Van Maanen’s Star, 138 
Vega, 13, 15, 30 
Veil Nebula (east), 145 
Veil Nebula (west), 144 
Virgo Cloud, 173 
VV Cephei, 140, 141 
VV Tauri, 71 

W Hydrae, 124 
W Ori, 127 

Whirlpool Galaxy, 167 
Wild Duck Cluster, 80 
Witch Head Nebula, 145 

X Cnc, 127 

V Ophiuchi, 121 

Zeta Orionis, 52 
Zeta Persei Association, 84 
Zubenelgenubi, 31 
Zubeneschamali, 18 

(3 Cet, 32 
p Cygni, 95 
(3 LMi, 32 
[3 Lyrae, 95 
[3 Vir, 32 
y 2 Vel, 141 
8 Cephei, 120 
8 Bootes, 95 
8 Lyrae, 95 
£ Ursae Majoris, 116 
l 2 Sco, 33 
T] Aquilae, 120 
"q Persei, 21 
0 Apodis, 34 
0 Orionis C, 27 
[i Canis Majoris, 94 
v^oo 33 
v 2 CMa 32 

A Ursae Majoris, 116 
o 2 Eridani, 40, 96, 138 
X Cygni, 124 
'P 1 Aurigae, 21